diff --git "a/batch_s000040.csv" "b/batch_s000040.csv"
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+++ "b/batch_s000040.csv"
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+source,target
+ For the three microquasars GRO J1655-40. XTE J1550-564 and GRS 1915+105 measurements of their masses have been made but there is so far no way to determine the angular momenta. so the test is necessarily restricted to their masses.," For the three microquasars GRO J1655-40, XTE J1550-564 and GRS 1915+105 measurements of their masses have been made but there is so far no way to determine the angular momenta, so the test is necessarily restricted to their masses."
+ Masses are usually derived from the determination of the orbit parameters in case of a binary., Masses are usually derived from the determination of the orbit parameters in case of a binary.
+ This holds for the three microquasars but also for the blaek hole AA* in the Galactic Center which 15 orbited by the S2 (SO-2) star (Schéddel et al..," This holds for the three microquasars but also for the black hole A* in the Galactic Center which is orbited by the S2 (S0-2) star (Schöddel et al.,"
+ 2002. Ghez et al..," 2002, Ghez et al.,"
+ 2003)., 2003).
+ Gezel et al. (, Genzel et al. (
+2003) have reported a QPO period of min observed in two new-infrared flares.,2003) have reported a QPO period of min observed in two near-infrared flares.
+ Aschenbach et al. (, Aschenbach et al. (
+20041) have claimed additionally quasi-periods around s. s. s and s. with the s period being consistent with the NIR period.,"2004) have claimed additionally quasi-periods around s, s, s and s, with the s period being consistent with the NIR period."
+ This set of quasi-periods was found in the power density spectra of one X-ray flare observed withChandra (Baganoff et al..," This set of quasi-periods was found in the power density spectra of one X-ray flare observed with (Baganoff et al.,"
+ 2001) and a second X-ray flare with (Porquet et al..," 2001) and a second X-ray flare with (Porquet et al.,"
+ 2003)., 2003).
+ Interesting in this context is that the frequencies corresponding to the latter three quasi-periods are close to a 3:2:1 ratio., Interesting in this context is that the frequencies corresponding to the latter three quasi-periods are close to a 3:2:1 ratio.
+ Enforeing such a ratio sequence in a best fit. it turns out that such a ratio is consistent with the measurements.," Enforcing such a ratio sequence in a best fit, it turns out that such a ratio is consistent with the measurements."
+ Accordingly. | include the Galactic Center black hole in this test and predict the mass of AA*.," Accordingly, I include the Galactic Center black hole in this test and predict the mass of A*."
+ The frequencies in question are 450 and 300 Hz (Strohmayer 2001a. Remillard et al.," The frequencies in question are 450 and 300 Hz (Strohmayer 2001a, Remillard et al."
+ 1999) with Oq- Ξ 450 Hz., 1999) with $\Omega\sb{V}$ = 450 Hz.
+ According to Equation 5. the predicted black hole mass is Ap;=6.760.13..., According to Equation \ref{eq:5} the predicted black hole mass is $M\sb{BH} = 6.76 \pm 0.1 M\sb{\odot}$ .
+ The relative uncertainty of the mass ts the same as that of the HFQPO measurements of z1.5%.., The relative uncertainty of the mass is the same as that of the HFQPO measurements of $\approx$.
+ Dynamical mass measurements have been reported by Orosz Bailyn (1997) with A/5;;=7.02d0.22M.. and more recently by Greene et al. (, Dynamical mass measurements have been reported by Orosz Bailyn (1997) with $M\sb{BH} = 7.02\pm 0.22 M\sb{\odot}$ and more recently by Greene et al. (
+2001) who obtained a mass range of pj;=5.8.G:8M,2001) who obtained a mass range of $M\sb{BH} = 5.8 - 6.8 M\sb{\odot}$.
+ The agreement between model prediction and observation is quite satisfactory., The agreement between model prediction and observation is quite satisfactory.
+ Frequencies of 276 and 184 Hz (Remillard et al..," Frequencies of 276 and 184 Hz (Remillard et al.,"
+ 2002) have been measured., 2002) have been measured.
+ With Οι: = 276 Hz Equation 5 predicts a black hole mass of Ap;=11.010.2ÀA7..., With $\Omega\sb{V}$ = 276 Hz Equation \ref{eq:5} predicts a black hole mass of $M\sb{BH} = 11.04 \pm 0.2 M\sb{\odot}$.
+ The relative uncertainty of the mass is again due to the accuracy of the HFQPO measurements of z2%.., The relative uncertainty of the mass is again due to the accuracy of the HFQPO measurements of $\approx$.
+ I»Dynamical mass measurements have been reported by Orosz et al. (, Dynamical mass measurements have been reported by Orosz et al. (
+2002a) with a +10 mass range of Mpj=S.I.11.6A/.. which is nicely matched by the predicted mass.,"2002a) with a $\pm$ $\sigma$ mass range of $M\sb{BH} = 8.4 - 11.6 M\sb{\odot}$, which is nicely matched by the predicted mass."
+ The relevant 3:2 pair at 168 and 113 Hz has only recently been found (Remillard et al..," The relevant 3:2 pair at 168 and 113 Hz has only recently been found (Remillard et al.,"
+ 2003. McClintock Remillard. 2004).," 2003, McClintock Remillard, 2004)."
+ With Ου: = 168 Hz Equation 5. predicts a black hole mass of lpj;=18.13+0.363..., With $\Omega\sb{V}$ = 168 Hz Equation \ref{eq:5} predicts a black hole mass of $M\sb{BH} = 18.13 \pm 0.36 M\sb{\odot}$.
+" The relative uncertainty of the mass is due the accuracy of the HFQPO measurements of z2%,..", The relative uncertainty of the mass is due the accuracy of the HFQPO measurements of $\approx$.
+ Dynamical mass measurements have been reported by Greiner at al. (, Dynamical mass measurements have been reported by Greiner at al. (
+"2001) with Afpy,,=110czLOAL.. and Harlaftis Greiner (2004) with 3p;=τιοτε11M...",2001) with $M\sb{BH} = 14.0 \pm 4.0 M\sb{\odot}$ and Harlaftis Greiner (2004) with $M\sb{BH} = 14.0 \pm 4.4 M\sb{\odot}$.
+ Also in this case the predicted mass matches the dyamically mass within the «1o range. but the measurements of this source illustrate nicely the potential uncertainty in cynamical mass measurements.," Also in this case the predicted mass matches the dynamically mass within the $\pm$ $\sigma$ range, but the measurements of this source illustrate nicely the potential uncertainty in dynamical mass measurements."
+ Harlaftis Greiner (2004) point out that the biggest uncertainty is related to the inclinatico of the orbit. which for GRS 19154105 is taken from the iclination of the associated jet assuming that the Jet axis ts orthogonal to the orbital plane of the binary.," Harlaftis Greiner (2004) point out that the biggest uncertainty is related to the inclination of the orbit, which for GRS 1915+105 is taken from the inclination of the associated jet assuming that the jet axis is orthogonal to the orbital plane of the binary."
+ If. for instance. the 1nclination angle is changed from the adopted to the estimated black hole mass is Mp;=16.945.0M. (Harlaftis Greiner. 2004).," If, for instance, the inclination angle is changed from the adopted $\sp{\deg}$ to $\sp{\deg}$ the estimated black hole mass is $M\sb{BH} = 16.9 \pm 5.9 M\sb{\odot}$ (Harlaftis Greiner, 2004)."
+ Actually. Kaiser et al. (," Actually, Kaiser et al. ("
+2004) suggest that the inclination angle is putting the best estimate for the dynamical black hole mass at 17.8 A... which is within the 316 measurement,"2004) suggest that the inclination angle is $\sp{\deg}$ putting the best estimate for the dynamical black hole mass at 17.8 $M\sb{\odot}$ , which is within the $\pm$ $\sigma$ measurement"
+The alternative technique we use to analyse the stellar populations of 3357 is to cross-correlate the galaxy spectrum with each SSP spectrum from the Vazdekisetal.(2010) model library.,The alternative technique we use to analyse the stellar populations of 357 is to cross-correlate the galaxy spectrum with each SSP spectrum from the \citet{2010MNRAS.404.1639V} model library.
+" For this purpose the model spectra are previously prepared to match the spectral range, velocity dispersion, and spectral resolution of the data."," For this purpose the model spectra are previously prepared to match the spectral range, velocity dispersion, and spectral resolution of the data."
+" Moreover, the galaxy and model spectra are rebinned logarithmically and normalized to remove the continua."," Moreover, the galaxy and model spectra are rebinned logarithmically and normalized to remove the continua."
+" In order to optimize the cross-correlation method for disentangling different stellar populations, it is necessary to adequately filter the spectra and multiply them by a cosine-bell-like function (Tonry&Davis1979).."," In order to optimize the cross-correlation method for disentangling different stellar populations, it is necessary to adequately filter the spectra and multiply them by a cosine-bell-like function \citep{1979AJ.....84.1511T}."
+" The importance of choosing a suitable filter lies in the possibility of getting rid of the noise in the spectrum; this purpose can be achieved by simply removing the largest wavenumbers, where the information about the shortest wavelength ranges is included."," The importance of choosing a suitable filter lies in the possibility of getting rid of the noise in the spectrum; this purpose can be achieved by simply removing the largest wavenumbers, where the information about the shortest wavelength ranges is included."
+" Therefore, the limit is imposed by the resolution of the data."," Therefore, the limit is imposed by the resolution of the data."
+" On the other hand, shorter wavenumbers contain information about wider spectral ranges, so possible residuals of the continuum removal due to errors in the flux calibration might also be filtered."," On the other hand, shorter wavenumbers contain information about wider spectral ranges, so possible residuals of the continuum removal due to errors in the flux calibration might also be filtered."
+" The drawback of this procedure is that it implies a power loss of the final cross-correlation function, specially when filtering short wavenumbers where most of the signal is included."," The drawback of this procedure is that it implies a power loss of the final cross-correlation function, specially when filtering short wavenumbers where most of the signal is included."
+" Apart from the filtering, it might be required to mask some regions in the original spectra, as is usually done in the full-spectrum fitting technique in the wavelength space."," Apart from the filtering, it might be required to mask some regions in the original spectra, as is usually done in the full-spectrum fitting technique in the wavelength space."
+ We tested different masks trying to avoid those features that are not well reproduced by the models due to mismatched abundance ratios: the CN in the blue spectral range and the Mg and Hf features in the red spectral range., We tested different masks trying to avoid those features that are not well reproduced by the models due to mismatched abundance ratios: the CN in the blue spectral range and the Mg and $\beta$ features in the red spectral range.
+" Again, these features contain most of the signal of the power spectrum so the choice of the masks has to be done very carefully in order to not lose most of the information."," Again, these features contain most of the signal of the power spectrum so the choice of the masks has to be done very carefully in order to not lose most of the information."
+" Finally, the peak height obtained for each correlation function is plotted against the model age and metallicity."," Finally, the peak height obtained for each correlation function is plotted against the model age and metallicity."
+ Since the cross-correlation profile reaches a higher value when object and template are more similar (getting a, Since the cross-correlation profile reaches a higher value when object and template are more similar (getting a
+model so that the 230GIIz flux is 3.4Jv.,model so that the $230\GHz$ flux is $3.4\Jy$.
+ The model can also fal by cooling too rapillv to be consistent wilh our neglect of cooling in the dynamical model., The model can also fail by cooling too rapidly to be consistent with our neglect of cooling in the dynamical model.
+ The Tables list a radiative efficiency 72μοιMc. where Lpoy is the bolometric Iuminositv. (1ntegrated over solid angle). and for comparison a thin disk efficiency. at the same ἄν.," The Tables list a radiative efficiency $\eta \equiv L_{\rm BOL}/\mdot c^2$, where $L_{\rm BOL}$ is the bolometric luminosity (integrated over solid angle), and for comparison a thin disk efficiency at the same $a_*$ ."
+" ranges between 5.4x10! for a,—0.5.Ti/T.=10 to 0.18 [or a,=0.98.75/7.1 (the thin disk efficiency for (he latter is 0.25)."," $\eta$ ranges between $5.4 \times 10^{-4}$ for $a_* = 0.5, \Trat = 10$ to $0.18$ for $a_* = 0.98, \Trat = 1$ (the thin disk efficiency for the latter is $0.25$ )."
+" Only in the e,=0.98. T/T.=1 model is the radiative efficiency. sufficiently high that cooling is likely to have a significant. effect on the GRMIID model."," Only in the $a_* = 0.98$, $\Trat = 1$ model is the radiative efficiency sufficiently high that cooling is likely to have a significant effect on the GRMHD model."
+ We will consider models with cooling in a [uture publication., We will consider models with cooling in a future publication.
+ Very [ew of the time averaged SEDs based on a single-temperature (7)/7.=1) models produce the correct o.," Very few of the time averaged SEDs based on a single-temperature $\Trat
+= 1$ ) models produce the correct $\alpha$."
+ The exception is edge-on lori (/=85dee) around [ast spinning black holes (model E and E)., The exception is edge-on tori $i=85\deg$ ) around fast spinning black holes (model E and F).
+ These models are ruled out. however. because they overproduce NIR and X-ray (lus.," These models are ruled out, however, because they overproduce NIR and X-ray flux."
+" For T;/T.=3 only the model with e,=0.94 seen al /=85dee agrees with the data.", For $\Trat=3$ only the model with $a_* = 0.94$ seen at $i=85\deg$ agrees with the data.
+ This is the best-bet model discussed im ??.., This is the best-bet model discussed in \ref{sec:4.1}.
+" For 7=85deg. models with spins below a,=0.94 (A. D and C) are ruled out by the inconsistent spectral slope. and models with hieher spins (E and F). allhoueh consistent with the observed a. overproduce the quiescent NIR and X-ray emission."," For $i=85\deg$, models with spins below $a_* = 0.94$ (A, B and C) are ruled out by the inconsistent spectral slope, and models with higher spins (E and F), although consistent with the observed $\alpha$, overproduce the quiescent NIR and X-ray emission."
+ All models with 7;/7.=3 observed al 7=5deg and 45cdeg are ruled out by the inconsistent a., All models with $\Trat=3$ observed at $i=5\deg$ and $45\deg$ are ruled out by the inconsistent $\alpha$.
+ For 7/7.=10. we find that all models with 7=85deg are ruled out by both a and violation of NIB. and X-ray. limits.," For $\Trat=10$, we find that all models with $i=85\deg$ are ruled out by both $\alpha$ and violation of NIR and X-ray limits."
+ For lower inclination angles (7=5deg.45 clea) a few models (E and F with /=5deg. and A. D. C. and D at ¢=15 clee) reproduce the observed a.," For lower inclination angles $i=5\deg,45\deg$ ) a few models (E and F with $i =
+5\deg$, and A, B, C, and D at $i = 45\deg$ ) reproduce the observed $\alpha$."
+ These models are consistent with N-ravs and NIB. limitations., These models are consistent with X-rays and NIR limitations.
+ Models E andF for 7=45deg are ruled oul by NIR ancl X-ray. limitations whereas models A. D. C and D for 7=5deg produce a which is too small.," Models E andF for $i=45\deg$ are ruled out by NIR and X-ray limitations whereas models A, B, C and D for $i=5\deg$ produce $\alpha$ which is too small."
+ What is the physical origin of (hese constraints?, What is the physical origin of these constraints?
+" The dependence on e, arises lareely because as a, increases (he inner edge of the clisk the ISCO — reaches deeper into the gravitational potential of the black hole. where the temperature and magnetic field strength are higher."," The dependence on $a_*$ arises largely because as $a_*$ increases the inner edge of the disk — the ISCO – reaches deeper into the gravitational potential of the black hole, where the temperature and magnetic field strength are higher."
+" In the disk mid-plaue. the temperature is a fraction of the virial temperature and scales with radius 0,xl/r. Dxαν. "," In the disk mid-plane, the temperature is a fraction of the virial temperature and scales with radius $\Theta_e \propto 1/r$. $B \propto 1/r$,"
+while the density. ~r. below the pressure maximunm.," while the density $\sim
+r$, below the pressure maximum."
+ Holding all else constant (which we do not: we hold the 230Gllz flix constant) this implies a higher peak [requency for svuchrotron emission. a constant Thomson depth (in our modelsthe Thomson depth at the ISCO is roughly constant. since the path length 1/~rgo but the density 7 rico). aud a larger energy boost per scattering 1z 1607. as can be seen in comparing models with different," Holding all else constant (which we do not: we hold the $230\GHz$ flux constant) this implies a higher peak frequency for synchrotron emission, a constant Thomson depth (in our modelsthe Thomson depth at the ISCO is roughly constant, since the path length $1/\sim \risco$ but the density $\sim \risco$ ), and a larger energy boost per scattering $A \approx 16 \Theta_e^2$ , as can be seen in comparing models with different"
+The spin composition of the proton in terms of its fundamental quark and eluon degrees of freedom is a central focus of proton structure.,The spin composition of the proton in terms of its fundamental quark and gluon degrees of freedom is a central focus of proton structure.
+ Whether the quark orbital angular momentum is zero or not is one of the kev points to solve this problem., Whether the quark orbital angular momentum is zero or not is one of the key points to solve this problem.
+ The importance of quark orbital angular momentum. whieh one might have taken to vanish in the ground state. has been evident since the work of Sehgal |1]..," The importance of quark orbital angular momentum, which one might have taken to vanish in the ground state, has been evident since the work of Sehgal \cite{seh74}."
+ The orbital angular momentum structure of the proton is, The orbital angular momentum structure of the proton is
+Ιω A μοντί2NLN],$\rho_{\rm local}$ and $N_{\rm HDF-N}^{(m<28.8)}$.
+ Results are plotted in− Fig.− 5.. ," Results are plotted in Fig. \ref{mod_halo}, ,"
+"where pi and Nieτί38,8]< are shown as unctons of ALfpsony. the absolute magnitude of the halo population objects."," where $\rho_{\rm local}$ and $N_{\rm HDF-N}^{(m<28.8)}$ are shown as functions of $M_{\rm F606W}$, the absolute magnitude of the halo population objects."
+ Tmiis figure implies (he existence of a large number of halo objects that should be present . ∐⊔∐↲∐↻∊↕∐⋯≸≟≼↲⋟, This figure implies the existence of a large number of halo objects that should be present in the HDF images.
+∖⊽⋅⋋≼≻↥↩⊔⋯↴↥↼∖↓↓↻↓⇄⋋↕⋟∖⇁↖≺↽↔↴↕⋅≼↲≀↧↴∩↲↕⋅⊔↥≀↧↴∐↴∿↴−≻↱≻∪↕∐≀↧↴↥≺∢≀↕⊔∖⊽≼↲⋟∖⊽⋅⊺∐↕⋟∖⊽↕⋅≼↲⋟∖⇁∏∐↕⋟∖⊽∐∪↥., Note that $N_{\rm HDF-N}^{(m<28.8)}$ is greater than $\sim$ 250 in all cases.
+ ⋅ - zmNN2pe . . :⋅ ⋅ compatible with the ILDE-N observations. where no obvious stars are present. except [or a ew 200h magnitude ones. (INawaler.1996:Floin.Gould.&Baheall1996).," This result is not compatible with the HDF-N observations, where no obvious stars are present, except for a few 20th magnitude ones, \citep{K96,F96}."
+". In conclusion. the observed excess in SBF-measured af, cannot be produced by objects belonging to the Milkv. Wax halo."," In conclusion, the observed excess in SBF-measured $\sigma_{\rm BG}^{2}$ cannot be produced by objects belonging to the Milky Way halo."
+ Otherwise a large number of resolved objects from this halo population would show up in the IIDE-N images. which is not tlie case.," Otherwise a large number of resolved objects from this halo population would show up in the HDF-N images, which is not the case."
+ If the observed ��]ς; excess cannot be produced by Milky Way halo objects. the only possibility is (hab it is caused by [aint galaxies.," If the observed $\sigma_{BG}^{2}$ excess cannot be produced by Milky Way halo objects, the only possibility is that it is caused by faint galaxies."
+" The large excess obtained in (he στ, with respect to the n(m)-estimated o7,; would imply an increase in the slope of n(m) al some magnitude fainter than my=28.8.", The large excess obtained in the SBF-measured $\sigma_{BG}^{2}$ with respect to the $n(m)$ -estimated $\sigma_{BG}^{2}$ would imply an increase in the slope of $n(m)$ at some magnitude fainter than $m_c=28.8$.
+" This slope can be computed by fitting the n(n)-estimated 07,; to our SBF-measured σ]ς; and taking the slope as a [ree parameter.", This slope can be computed by fitting the $n(m)$ -estimated $\sigma_{BG}^{2}$ to our SBF-measured $\sigma_{BG}^{2}$ and taking the slope as a free parameter.
+" If it is assumed that the slope change occurs αἱ m,=28.8 for all filters. the resulting slopes for the [ainter range are 5?=0.60. 0.44. ancl 0.54 for Bysy. Ving. ancl {κι respectively."," If it is assumed that the slope change occurs at $m_c = 28.8$ for all filters, the resulting slopes for the fainter range are $\gamma=0.60$, $0.44$, and $0.54$ for $B_{450}$, $V_{606}$, and $I_{814}$, respectively."
+" These slopes would be valid up to σου=34.4. Ving=31.9. ancl fy)=32.5 al least since the contribution of fainter magnitudes to o7; becomes smaller than the uncertainties in the SBF-measured 67,; results."," These slopes would be valid up to $B_{450}=34.4$, $V_{606}=31.9$, and $I_{814}=32.5$ at least since the contribution of fainter magnitudes to $\sigma_{BG}^{2}$ becomes smaller than the uncertainties in the SBF-measured $\sigma_{BG}^{2}$ results."
+ Hf the slope change were to occur at a magnitude fainter than 28.8. it would result in a steeper vr).," If the slope change were to occur at a magnitude fainter than 28.8, it would result in a steeper $n(m)$."
+ In anv case. such bie changes in the slope of n(m) seem unrealistic.," In any case, such big changes in the slope of $n(m)$ seem unrealistic."
+ In our opinion. this possibility should be rejected.," In our opinion, this possibility should be rejected."
+ As a consequence. it must be concluded that the Williamsetal.(1996) data are incomplete.," As a consequence, it must be concluded that the \citet{W96} data are incomplete."
+" Asstuning that the Metealleοἱal.(2001). differential number counts are correct. the SDF-measured o5, results listed in Table 7 and the n(m)-estimated oj; values obtained using the Metcalfeetal.(2001). data. listed in Table 3.. can be compared."," Assuming that the \citet{Met01} differential number counts are correct, the SBF-measured $\sigma_{\rm BG}^2$ results listed in Table \ref{t-results-bg} and the $n(m)$ -estimated $\sigma_{\rm BG}^2$ values obtained using the \citet{Met01} data, listed in Table \ref{t-sigma}, can be compared."
+ It can be seen that the SBF-measured and n(m)-estimated σος coincide withinthe error bars for the FSl4W filter. and is very similar for the FASOW filter.," It can be seen that the SBF-measured and $n(m)$ -estimated $\sigma_{\rm BG}^2$ coincide withinthe error bars for the F814W filter, and is very similar for the F450W filter."
+ Only in the filter F606W some differences arise., Only in the filter F606W some differences arise.
+ This implies that extrapolation of the Metcalfe n(n) bunction to magnitudes fainter than 28.8 accounts almost entirely for the measured, This implies that extrapolation of the \citet{Met01} $n(m)$ function to magnitudes fainter than 28.8 accounts almost entirely for the measured
+ealaxiesnupliceatious.andgalaxyIufact.if Dark Matter (DAD aunibhilatious in the halo of clusters have crucial astrophysical DM.,"}\tikzmark{mainBodyEnd0}
+\date{Received 16 October 2006 / Accepted 25 January 2007 }
+
+\authorrunning {S. Colafrancesco et al.}
+
+\titlerunning {SZ effect from 1ES0657-556 cluster}
+
+\abstract
+ % contex\tikzmark{mainBodyStart1}xt\tikzmark{mainBodyEnd1} \tikzmark{mainBodyStart2}heading\tikzmark{mainBodyEnd2} \tikzmark{mainBodyStart3}(optional)\tikzmark{mainBodyEnd3}
+ %{} leave it empty if necessary
+ {}\tikzmark{mainBodyStart4}}\tikzmark{mainBodyEnd4}
+ % aims heading (mandatory)
+ {The cluster \es is an ideal astrophysical laboratory to study the
+ distribution and the nature of Dark Matter because this last component is spatially
+ separated from the intracluster gas.
+ We show that microwave observations can provide crucial probes of Dark Matter in this
+ system.}\tikzmark{mainBodyStart5}}\tikzmark{mainBodyEnd5}
+ % methods heading (mandatory)
+ {We calculate the expected SZ effect from Dark Matter annihilation in the main mass
+ concentrations of the cluster 1ES0657-556, and we estimate the sources of contamination,
+ confusion and bias to asses its significance.}\tikzmark{mainBodyStart6}}\tikzmark{mainBodyEnd6}
+ % results heading (mandatory)
+ {We find that SZ observations at $\nu \approx 223$ GHz can resolve both spatially and
+ spectrally the SZ$_{DM}$ signal and isolate it from the other SZ signals,
+ and mainly from the thermal SZ effect which is null at $ \nu \sim 220-223$ GHz
+ for the case of \es. We conclude that SZ observations with $\simlt$ arcmin resolution
+ and $\simlt \mu$K sensitivity of \es are crucial, and maybe unique, to find direct
+ astrophysical probes of the existence and of the nature of Dark Matter,
+ or to set strong experimental limits.}\tikzmark{mainBodyStart7}}\tikzmark{mainBodyEnd7}
+ % conclusions heading (optional), leave it empty if necessary
+ {}\tikzmark{mainBodyStart8}} Dark Matter (DM) annihilations in the halo of galaxies and galaxy clusters have crucial astrophysical implications."
+ is constituted by weakly interacting inassive particles (for which the leaciug candidate is the liehtest supersviuinetrie particle. plausibly the ucutralino X). their annihilation produces secondary particles (e.g. neutral aud charged pious. secondary electrons and protons. neutrinos) that eive rise to various astrophysical signals.," In fact, if DM is constituted by weakly interacting massive particles (for which the leading candidate is the lightest supersymmetric particle, plausibly the neutralino $\chi$ ), their annihilation produces secondary particles (e.g., neutral and charged pions, secondary electrons and protons, neutrinos) that give rise to various astrophysical signals."
+ These are. amoung others. observable fluxes of positrous. antiprotous. eamunua raves. neutrinos. as well as signals due to secondary clectrous which COVOT the whole eau.," These are, among others, observable fluxes of positrons, antiprotons, gamma rays, neutrinos, as well as signals due to secondary electrons which cover the whole e.m."
+" spect (sec, e.g.. Colafrancesco ct al."," spectrum (see, e.g., Colafrancesco et al."
+ 2006 for details): svuchrotron radio cussion (in the iutra-cluster maeuctic field). brenisstrahnluus cussion Gf there ds co-spatial iutra-cluster eas). inverse Compton cussion due to the upescatteriue of CAIB photous aud lence a specific SZ effect (as first noticed aud derived by Colafraucesco 20080.," 2006 for details): synchrotron radio emission (in the intra-cluster magnetic field), bremsstrahlung emission (if there is co-spatial intra-cluster gas), inverse Compton emission due to the up-scattering of CMB photons and hence a specific SZ effect (as first noticed and derived by Colafrancesco 2004)."
+ The spatial and spectral intensity of the astrophysical siguals coming frou xy annihlilation is expected. however. to be confused or even overcome bv other astrophysical sicnals originating from the intracluster (IC) eas aud/or from). the relativistic plasmas present iu the cluster atinospheres. especially when all these componcuts are co-spatially distributed with he DM componeut.," The spatial and spectral intensity of the astrophysical signals coming from $\chi \chi$ annihilation is expected, however, to be confused or even overcome by other astrophysical signals originating from the intracluster (IC) gas and/or from the relativistic plasmas present in the cluster atmospheres, especially when all these components are co-spatially distributed with the DM component."
+ This situation occurs in nost galaxw clusters (see ciscussion by Colatrancesco et al., This situation occurs in most galaxy clusters (see discussion by Colafrancesco et al.
+ 2006 or the case of An ideal system to detect DALI auuihilation siguals would. therefore. be a clhuster with a clear spatial separation between the various matter components.," 2006 for the case of An ideal system to detect DM annihilation signals would, therefore, be a cluster with a clear spatial separation between the various matter components."
+ This is. indeed. the case of the cluster 1ES0657-556 where the spatial distribution of DAL is clearly offset παν," This is, indeed, the case of the cluster 1ES0657-556 where the spatial distribution of DM is clearly offset w.r.t."
+ that of the IC eas (Clowe et al., that of the IC gas (Clowe et al.
+ 2006)., 2006).
+ The two barvonic clumps of hot eas cuit A-ravs by (thermal) broenisstrahluug. as Observer by Chaudra (AMarkevitch et al.," The two baryonic clumps of hot gas emit X-rays by (thermal) bremsstrahlung, as observed by Chandra (Markevitch et al."
+ 2002. 2001).," 2002, 2004)."
+ The shock observed in the westeruanost region of the cluster (Markevitch et al., The shock observed in the western-most region of the cluster (Markevitch et al.
+ 2002) wieght be the site of high euergv chussion from particles accelerated at the shock., 2002) might be the site of high energy emission from particles accelerated at the shock.
+ Ward X-rav elussion from the direction of 1ES0657-556 has been mareiually detected by Rossi-XTE (Petrosian et al., Hard X-ray emission from the direction of 1ES0657-556 has been marginally detected by Rossi-XTE (Petrosian et al.
+ 2006) but its augular resolution is uot sufficient to eive iv information on the spatial distribution of this cussion., 2006) but its angular resolution is not sufficient to give any information on the spatial distribution of this emission.
+ No eanmua-rav Cluission has been detected from this svsteus with EGRET., No gamma-ray emission has been detected from this system with EGRET.
+ The extended radio halo associated to this cluster (Liane ¢ al., The extended radio halo associated to this cluster (Liang et al.
+ 2000) has a surface brightuess slieltly clongated along the direction of the two X-ray chumps. but without clear signatures ofradio-briehtuess euliancemeuts at the DM. chunp locations.," 2000) has a surface brightness slightly elongated along the direction of the two X-ray clumps, but without clear signatures of radio-brightness enhancements at the DM clump locations."
+ Finally. the SZ maps of 1ES0657-556 obtained with ACBAR (with ~L5 arciniu FEWIIM resolution. Gomez et al.," Finally, the SZ maps of 1ES0657-556 obtained with ACBAR (with $\sim 4.5$ arcmin FWHM resolution, Gomez et al."
+ 2003) are quite smooth iid regular wi1 no evidence of chhancement at both X-rav and/or DM clump locations., 2003) are quite smooth and regular with no evidence of enhancement at both X-ray and/or DM clump locations.
+ Iu this Letter. we will compute the specific feature of the SZ effect. (horefator SZE) produced by DM annihilation. SZpy. iu the cluster aud we will show that it is possible to detect such SZpay signal with a specific observational strategy.," In this Letter, we will compute the specific feature of the SZ effect (herefater SZE) produced by DM annihilation, $_{\rm DM}$, in the cluster and we will show that it is possible to detect such $_{\rm DM}$ signal with a specific observational strategy."
+" The relevaut physical quautities are calculated using 77,= Tülans + banda flat. ACDAL (Q,,=0.3.04 0.7) cosmological model."," The relevant physical quantities are calculated using $H_0 = 70$ km $^{-1}$ $^{-1}$ and a flat, $\Lambda$ CDM $\Omega_{\rm m} = 0.3, \Omega_{\Lambda}=0.7$ ) cosmological model."
+" The various SZ signals expected from the subsvsteuis of the cluster are: 1) the SZpa, effect. which is expected to be located at the two DM chips: ii) the thermal SZ effect (SZin) which is expected to be located at the two N-rav chuups."," The various SZ signals expected from the subsystems of the cluster are: i) the $_{\rm DM}$ effect, which is expected to be located at the two DM clumps; ii) the thermal SZ effect $_{\rm th}$ ) which is expected to be located at the two X-ray clumps."
+ We will compute in the following these two sources of SZE aud we will also discuss the possible sources of coutamunation. bias aud confusion.," We will compute in the following these two sources of SZE and we will also discuss the possible sources of contamination, bias and confusion."
+ The seueral expression for the SZE which is valid in the Thotsou Πιτ for a eeneric electron population in the relativistic hit aud includes also the effects of unultiple scatteriues and the combination with other electron population iu the cluster atmospheres hay. been derived by Colafraucesco et al. (, The general expression for the SZE which is valid in the Thomson limit for a generic electron population in the relativistic limit and includes also the effects of multiple scatterings and the combination with other electron population in the cluster atmospheres has been derived by Colafrancesco et al. (
+2003).,2003).
+ This approac[um is the one that will be used for the derivation of the SZpyy effect induce by the secondary electrons produced by \ αλαπο (see the original derivation by Colafraucesco, This approach is the one that will be used for the derivation of the $_{\rm DM}$ effect induced by the secondary electrons produced by $\chi \chi$ annihilation (see the original derivation by Colafrancesco
+between different planet masses when comparing actual measurements to evolution nary models.,between different planet masses when comparing actual measurements to evolution nary models.
+" We see two major effects: (i) the photometric performance clearly depends on wavelength, and (ii) there are two different regimes depending on the position compared to the AO control radius."," We see two major effects: (i) the photometric performance clearly depends on wavelength, and (ii) there are two different regimes depending on the position compared to the AO control radius."
+" The first effect is directly related to the chromaticity of the PSF: in speckle-limited regime the noise attenuation is almost constant with angular separation compared to the coronagraphic profile, and the level of the coronagraphic profile linearly depends on wavelength."," The first effect is directly related to the chromaticity of the PSF: in speckle-limited regime the noise attenuation is almost constant with angular separation compared to the coronagraphic profile, and the level of the coronagraphic profile linearly depends on wavelength."
+ The second effect is related to the AO correction inside the control radius., The second effect is related to the AO correction inside the control radius.
+" Inside that region we see a stabilization of the performance: 0.2 mag photometric precision can be reached up to contrast of 10 to 11 mag (10* to 4x 10*) from to the AO control radius, which extends from in Y band to in K band."," Inside that region we see a stabilization of the performance: 0.2 mag photometric precision can be reached up to contrast of 10 to 11 mag $10^{-4}$ to $4 \times 10^{-4}$ ) from to the AO control radius, which extends from in Y band to in K band."
+" Outside of the AO control radius, the photometric performance increases almost linearly with angular separation at all wavelengths to reach contrast values of 14 to 15 mag (2.5x10 to 10 9) around2."," Outside of the AO control radius, the photometric performance increases almost linearly with angular separation at all wavelengths to reach contrast values of 14 to 15 mag $2.5 \times 10^{-6}$ to $10^{-6}$ ) around."
+"0"".. These numbers are given in the context? of our simulated test case, but the general effects should be similar for any data obtained with high contrast coronagraphic imagers."," These numbers are given in the context of our simulated test case, but the general effects should be similar for any data obtained with high contrast coronagraphic imagers."
+" Similarly to the noise level, using the SDI+ADI data analysis method improves the photometric accuracy."," Similarly to the noise level, using the SDI+ADI data analysis method improves the photometric accuracy."
+" However, using SDI+ADI will only provide an estimation of the differential flux of the planet between the 2 filters, contrary to ADI which provides an absolute measurement."," However, using SDI+ADI will only provide an estimation of the differential flux of the planet between the 2 filters, contrary to ADI which provides an absolute measurement."
+" To preserve the planet differential flux, the amplitude correction factor usually applied for SDI in the subtraction is taken equal to a fixed value of 1."," To preserve the planet differential flux, the amplitude correction factor usually applied for SDI in the subtraction is taken equal to a fixed value of 1."
+" The photometric error estimated with SDI+ADI follows the same variations as for ADI, but at higher contrast values."," The photometric error estimated with SDI+ADI follows the same variations as for ADI, but at higher contrast values."
+" Figure 2,, right, illustrates the photometric performance as a function of wavelength and angular separation in SDIJ-ADI."," Figure \ref{fig:flux_error_summary}, right, illustrates the photometric performance as a function of wavelength and angular separation in SDI+ADI."
+" The trends are similar to ADI alone, but the chromatic effect is less significant because the PSF chromaticity has been mitigated by the SDI part of the analysis."," The trends are similar to ADI alone, but the chromatic effect is less significant because the PSF chromaticity has been mitigated by the SDI part of the analysis."
+" Compared to ADI alone, the contrast values at which a 0.2 mag photometric error is reached are 1.5 to 2.5 mag higher."," Compared to ADI alone, the contrast values at which a 0.2 mag photometric error is reached are 1.5 to 2.5 mag higher."
+" At shorter wavelengths, in Y2Y3 filters, performances at separations larger than decrease."," At shorter wavelengths, in Y2Y3 filters, performances at separations larger than decrease."
+" This effect is related to the size of the aperture for photometry which is very small in Y band (4 pixels in diameter), and to the field rotation which has strong effect on encircled energy at separations larger thana in Y band."," This effect is related to the size of the aperture for photometry which is very small in Y band (4 pixels in diameter), and to the field rotation which has a strong effect on encircled energy at separations larger than in Y band."
+ Considering shorter exposures for individual images where the field rotation is negligible would decrease the photometric errors in that particular case., Considering shorter exposures for individual images where the field rotation is negligible would decrease the photometric errors in that particular case.
+ We hereafter combine the photometric accuracy obtained in ADI and SDI+ADI to define empirical photometric error curves for each filter pair as a function of contrast., We hereafter combine the photometric accuracy obtained in ADI and SDI+ADI to define empirical photometric error curves for each filter pair as a function of contrast.
+" The photometric error curves as a function of contrast at each angular separation have been fitted with the empirically defined function: where phote is the photometric error, c the contrast and (p1,p2,pa) the fitted parameters."," The photometric error curves as a function of contrast at each angular separation have been fitted with the empirically defined function: where $\mathrm{phot_{err}}$ is the photometric error, $c$ the contrast and $(p_1, p_2, p_3)$ the fitted parameters."
+ This function approaches the measured points with a precision of ~1%.., This function approaches the measured points with a precision of $\sim$.
+ The fitting has been performed for ADI and SDI+ADI., The fitting has been performed for ADI and SDI+ADI.
+" To take into account the scattering of the error with the planet position in the images, different cases have been considered at each separation, corresponding to the 3 different simulated planet positions: a standard case with an average photometric error, an optimal case corresponding to the lowest estimation of the error and a pessimistic case corresponding to the upper estimation of the error."," To take into account the scattering of the error with the planet position in the images, different cases have been considered at each separation, corresponding to the 3 different simulated planet positions: a standard case with an average photometric error, an optimal case corresponding to the lowest estimation of the error and a pessimistic case corresponding to the upper estimation of the error."
+ These empirical photometric errors are plotted in Fig., These empirical photometric errors are plotted in Fig.
+ 3 for the 4 simulated filter pairs., \ref{fig:error_curves_phot_all} for the 4 simulated filter pairs.
+ The amplitude of the error bars is defined by the optimal and pessimistic error curves described above., The amplitude of the error bars is defined by the optimal and pessimistic error curves described above.
+" We assume that the photometric error in ADI is the same in the two filters of a pair, which is legitimate given the amplitude of the error bars."," We assume that the photometric error in ADI is the same in the two filters of a pair, which is legitimate given the amplitude of the error bars."
+ These empirical error curves lie in the same range as the expected photometric accuracy of other data analysis methods developed within the SPHERE consortium by ? and ?.., These empirical error curves lie in the same range as the expected photometric accuracy of other data analysis methods developed within the SPHERE consortium by \citet{mugnier2008} and \citet{smith2009}.
+ Table 3 gives for each filter pair and each angular separation the contrast value at which the photometric error in ADI becomes lower than the differential photometric error in SDI+ADI., Table \ref{tab:methods_limits} gives for each filter pair and each angular separation the contrast value at which the photometric error in ADI becomes lower than the differential photometric error in SDI+ADI.
+ These values give the contrast at which it becomes more interesting in terms of photometric error to obtain a differential flux estimation., These values give the contrast at which it becomes more interesting in terms of photometric error to obtain a differential flux estimation.
+" As explained in Sect. 4.4,,"," As explained in Sect. \ref{sec:photometric_accuracy_sdi_adi},"
+ aperture photometry in Y band is extremely sensitive to errors introduced by the position of the aperture or the field rotation because the aperture is very small., aperture photometry in Y band is extremely sensitive to errors introduced by the position of the aperture or the field rotation because the aperture is very small.
+ This is why in Y2Y3 pair at there is no contrast limit between ADI and SDI+ADI: for that particular case the flux estimation error is slightly better in ADI than SDI+ADI., This is why in Y2Y3 pair at there is no contrast limit between ADI and SDI+ADI: for that particular case the flux estimation error is slightly better in ADI than SDI+ADI.
+" In this section we evaluate characterization capabilities of IRDIS in imaging mode, ie. how well the physical parameters and of the planets can be estimated from photometric measurements in different spectral bands."," In this section we evaluate characterization capabilities of IRDIS in imaging mode, i.e. how well the physical parameters and of the planets can be estimated from photometric measurements in different spectral bands."
+" To estimate the characterization capabilities of IRDIS, we performed a new simulation using as input the ὅ-σ detection limits obtained from Sect."," To estimate the characterization capabilities of IRDIS, we performed a new simulation using as input the $\sigma$ detection limits obtained from Sect."
+ 4.1 and the empirical error curves obtained in Sect. 4.5.., \ref{sec:noise_level} and the empirical error curves obtained in Sect. \ref{sec:empirical_photometric_accuracy}.
+ The goal of the simulation was to test the efficiency of all filter pair sequences for characterization, The goal of the simulation was to test the efficiency of all filter pair sequences for characterization
+frequency shift was affected.,frequency shift was affected.
+ This resulted in an estimated stray-light temperature of ~35 K which corresponds to ~4 pW of parasitic power per pixel at 150 GHz., This resulted in an estimated stray-light temperature of $\sim 35$ K which corresponds to $\sim 4$ pW of parasitic power per pixel at 150 GHz.
+ The unwanted radiation has thus been reduced by more than a factor of two compared with the first generation NIKA and is now comparable to the best sky conditions at Pico Veleta., The unwanted radiation has thus been reduced by more than a factor of two compared with the first generation NIKA and is now comparable to the best sky conditions at Pico Veleta.
+ The dual-band NIKA run took place in October 2010., The dual-band NIKA run took place in October 2010.
+ The instrument was installed in the receiver cabin of the IRAM 30-meter telescope at Pico Veleta. Spain. and operated remotely from the control room.," The instrument was installed in the receiver cabin of the IRAM 30-meter telescope at Pico Veleta, Spain, and operated remotely from the control room."
+ The cool-down of the instrument was also performed remotely. taking approximately 18 hours to reach the operating temperature of 70 mK. Astronomical data from the two arrays are reduced with dedicated software.," The cool-down of the instrument was also performed remotely, taking approximately 18 hours to reach the operating temperature of 70 mK. Astronomical data from the two arrays are reduced off-line with dedicated software."
+ The raw data (/.Q) are converted to complex phase angle using the closest previous KID calibration., The raw data $I$ $Q$ ) are converted to complex phase angle using the closest previous KID calibration.
+ Then a conversion to an equivalent frequency shift is done with the same calibration using the derivative of the frequency with the complex phase at the zero phase. as described in4.," Then a conversion to an equivalent frequency shift is done with the same calibration using the derivative of the frequency with the complex phase at the zero phase, as described in."
+. Data are thus internally converted to frequencies which are assumed to be linear with the absorbed photon counts. as in equation(2).," Data are thus internally converted to frequencies which are assumed to be linear with the absorbed photon counts, as in equation."
+. After opacity correction. and using Mars as the primary calibrator. we obtain that the overall median gain is of 14 mJy/beam/Hz and 9 mJy/beam/Hz for the I.4 and 2 mm (220 GHz and 150 GHz) channels. with a 30% dispersion.," After opacity correction, and using Mars as the primary calibrator, we obtain that the overall median gain is of 14 mJy/beam/Hz and 9 mJy/beam/Hz for the 1.4 and 2 mm (220 GHz and 150 GHz) channels, with a $30\%$ dispersion."
+ The focal plane geometry of each array ts measured by using scanning maps of planets (see 7))., The focal plane geometry of each array is measured by using scanning maps of planets (see ).
+ The fitted focal plane geometry is found by matching the pixel position in the array as measured on the wafer to the measured position on planets. by optimizing a simple set of parameters: a center. a tilt angle and a sealing expressed in areseconds/mm.," The fitted focal plane geometry is found by matching the pixel position in the array as measured on the wafer to the measured position on planets, by optimizing a simple set of parameters: a center, a tilt angle and a scaling expressed in arcseconds/mm."
+ Most detectors are within less than 2 arcseconds of their expected position., Most detectors are within less than 2 arcseconds of their expected position.
+ The beam width is also found from planet measurements., The beam width is also found from planet measurements.
+ Typically the FWHM is 12.4 and 16.7 areseconds for the two arrays (1.4 and 2 mm respectively. see figure 8) with a dispersion of | aresecond.," Typically the FWHM is 12.4 and 16.7 arcseconds for the two arrays (1.4 and 2 mm respectively, see figure 8) with a dispersion of 1 arcsecond."
+ This is close to the diffraction limit for the 2 mm array., This is close to the diffraction limit for the 2 mm array.
+ Pixelisation, Pixelisation
+magnetar flares can point to the location where a long GRB (the time it takes for B—[0' G fields toY.,magnetar flares can point to the location where a long GRB (the time it takes for $B\sim 10^{16}$ G fields to.
+. Here. I show that the GRB afterglow emission should be still detectable in the radio when the superflare takes place.," Here, I show that the GRB afterglow emission should be still detectable in the radio when the superflare takes place."
+ GRB afterglows can be followed in the radio wavelengths for years after the burst., GRB afterglows can be followed in the radio wavelengths for years after the burst.
+ GRB 030329 is an intrinsically typical long GRB that took place particularly nearby at z=0.1685 tor luminosity distance of dg;=800 Mpe for standard cosmology: Greiner et al., GRB 030329 is an intrinsically typical long GRB that took place particularly nearby at $z=0.1685$ (or luminosity distance of $d_L=800$ Mpc for standard cosmology; Greiner et al.
+ 2003)., 2003).
+ Its radio afterglow remains fairly bright (at the mJy level) years after the burst and the blastwave is resolved (e.g. Berger et al., Its radio afterglow remains fairly bright (at the mJy level) years after the burst and the blastwave is resolved (e.g. Berger et al.
+ 2003: Tavlor et al., 2003; Taylor et al.
+ 2004: Resmi et al., 2004; Resmi et al.
+ 2005: Frail et al., 2005; Frail et al.
+ 2005: Pihlstrómm et al., 2005; Pihlströmm et al.
+ 2007: van der Horst et al., 2007; van der Horst et al.
+ 2008)., 2008).
+ Because of the slow decline in flux. the afterglow is expected to be observable over the next decade in the GHz range and beresolved ~7 years after the burst (Pihlstrómm et al.," Because of the slow decline in flux, the afterglow is expected to be observable over the next decade in the $GHz$ range and be $\sim$ 7 years after the burst (Pihlströmm et al."
+ 2007)., 2007).
+ With the Low Frequency Array (LOFAR) the afterglow of GRB 030329 can be detected for several decades (van der Horst et al., With the Low Frequency Array ) the afterglow of GRB 030329 can be detected for several decades (van der Horst et al.
+ 2008)., 2008).
+ The afterglow emission of a GRB similar to that of 030329 located at a distance d;~250 Mpe will be ~10 times more bright and with the radio image a factor of ~2.6 larger., The afterglow emission of a GRB similar to that of 030329 located at a distance $d_L\sim 250$ Mpc will be $\sim 10$ times more bright and with the radio image a factor of $\sim 2.6$ larger.
+ Such an afterglow emission can be detected and resolved for hundred cor hundreds) of years after the burst., Such an afterglow emission can be detected and resolved for hundred (or hundreds) of years after the burst.
+ Two-dimensional relativistic hydrodynamical simulations (Zhang and MacFadyen 2009) indicate that the GRB blast reaches a distance of ~3 peat ~100 years which corresponds to a source of angular size of ~2.7 mas (for a corresponding angular distance of ας~224 Mpc) and flux density of ~0.1 mJy (at ~ IGHz) allowing for the morphological study of the blastwave with high-sensitivity Very Long Baseline Interferometry (VLBI) observations similar to those reported in Pihlstrómm et al. (, Two-dimensional relativistic hydrodynamical simulations (Zhang and MacFadyen 2009) indicate that the GRB blast reaches a distance of $\sim 3$ pc at $\sim 100$ years which corresponds to a source of angular size of $\sim 2.7$ mas (for a corresponding angular distance of $d_A\sim 224$ Mpc) and flux density of $\sim$ 0.1 mJy (at $\sim 1$ GHz) allowing for the morphological study of the blastwave with high-sensitivity Very Long Baseline Interferometry (VLBI) observations similar to those reported in Pihlströmm et al. (
+2007).,2007).
+ According to the same simulations. the decelerating GRB blastwave is morphologically very different from a supernova remnant for the first ~200 years allowing for the distinction between the two types of explosions.," According to the same simulations, the decelerating GRB blastwave is morphologically very different from a supernova remnant for the first $\sim 200$ years allowing for the distinction between the two types of explosions."
+ For radio follow-ups to be possible. a good enough localization of the superflare is needed.," For radio follow-ups to be possible, a good enough localization of the superflare is needed."
+ Such localization can be provided with the Burst Alert Telescope (BAYT) detector onSW/FT., Such localization can be provided with the Burst Alert Telescope ) detector on.
+ The rate at whichSWIFT detects GRBs is ~ 1/3 of that ofBATSE mainly because of its smaller field of view., The rate at which detects GRBs is $\sim$ 1/3 of that of mainly because of its smaller field of view.
+ I. thus. estimate that detects Ry/3-.. superflares per year.," I, thus, estimate that detects ${\dot R}_{\rm sf}/3\sim 1fL_{49}^{3/2}$ superflares per year."
+ detection rate of flares is a factor of «2.5 higher but the Glast Burst Monitor (GBAT) lacks the localization needed for a radio follow-up., detection rate of flares is a factor of $\sim$ 2.5 higher but the Glast Burst Monitor ) lacks the localization needed for a radio follow-up.
+ The pulsating tail that is expected to follow the superflare may. in some cases. be powerful enough to be observed withXA hundreds of seconds after the event.," The pulsating tail that is expected to follow the superflare may, in some cases, be powerful enough to be observed with hundreds of seconds after the event."
+ Although the pulsating tails that follow bright SGR flares of Galactic magnetars for ~200— 400s have Lay107 ergs/s (Mereghetti 2008). the strong magnetic field of the GRB magnetar can confine ~100 times more energy in the magnetosphere of the neutron star resulting in far brighter ray tails.," Although the pulsating tails that follow bright SGR flares of Galactic magnetars for $\sim 200-400$ s have $L_{\rm tail}\sim 10^{42}$ ergs/s (Mereghetti 2008), the strong magnetic field of the GRB magnetar can confine $\sim 100$ times more energy in the magnetosphere of the neutron star resulting in far brighter X-ray tails."
+ It is furthermore possible that the superflare has a strong enough “afterglow” of its own that allows for X-ray (tor longer wavelength) detection and accurate localization shortly after the burst (Eichler 2002)., It is furthermore possible that the superflare has a strong enough “afterglow” of its own that allows for X-ray (or longer wavelength) detection and accurate localization shortly after the burst (Eichler 2002).
+ If GRB-magnetars exist. their magnetic field should decay on a time-scale of a few hundred years possibly producing SGR-like flares with peak luminosities of .0 ergs/s. A few of these flares per year should have been detected by BATSE out to d;~250 Mpe classified as short-duration GRBs.," If GRB-magnetars exist, their magnetic field should decay on a time-scale of a few hundred years possibly producing SGR-like flares with peak luminosities of $\sim 10^{49}$ ergs/s. A few of these flares per year should have been detected by BATSE out to $d_L\sim 250$ Mpc classified as short-duration GRBs."
+ Such superflares can be detected withSWIFT at a rate of about one per year., Such superflares can be detected with at a rate of about one per year.
+ The host galaxy of the flare should be typical of those of long-duration GRBs., The host galaxy of the flare should be typical of those of long-duration GRBs.
+ High sensitivity radio observations at the location of the flare can resolve a~ [00-year-old blastwave result of the interaction of the GRB jets with the circumburst medium., High sensitivity radio observations at the location of the flare can resolve a $\sim 100$ -year-old blastwave result of the interaction of the GRB jets with the circumburst medium.
+ This detection can prove that GRBs are connected to the birth of magnetars., This detection can prove that GRBs are connected to the birth of magnetars.
+ I thank Brian Metzger and Dmitri Uzdensky for stimulating discussions during the preparation of the manuscript., I thank Brian Metzger and Dmitri Uzdensky for stimulating discussions during the preparation of the manuscript.
+ I acknowledge support from the Lyman Spitzer. Jr. Fellowship awarded by the Department of Astrophysical Sciences at Princeton University.," I acknowledge support from the Lyman Spitzer, Jr. Fellowship awarded by the Department of Astrophysical Sciences at Princeton University."
+"(Tor our fiducial choices),",(for our fiducial choices).
+" At the small apertures we will consider. the diffraction limit is larger than the seeing. so it is possible to make (he PSF much smaller (han a pixel. ρω<06,."," At the small apertures we will consider, the diffraction limit is larger than the seeing, so it is possible to make the PSF much smaller than a pixel, $\theta_{\rm PSF}\ll \theta_p$."
+ This would have the advantage of reducing sky. noise and is a useful approach when it is possible to always center the telescope at the same fiekl position as is the case lor “point and stare” experiments., This would have the advantage of reducing sky noise and is a useful approach when it is possible to always center the telescope at the same field position as is the case for “point and stare” experiments.
+ Hlowever. for an all-sky survey. which eveles through maux fields. such precision repeat pointing is extremely difficult.," However, for an all-sky survey, which cycles through many fields, such precision repeat pointing is extremely difficult."
+ Without it. precision photometry is impossible unless (he sub-pixel response of the CCD is mapped out in detail.," Without it, precision photometry is impossible unless the sub-pixel response of the CCD is mapped out in detail."
+ We therefore adopt a Nyquist-sampled PSF. for which the sky noise is approximately that falling on da~13 pixels.," We therefore adopt a Nyquist-sampled PSF, for which the sky noise is approximately that falling on $4\pi\sim 13$ pixels."
+ Our overall consideration for telescope design must (ake into account three factors., Our overall consideration for telescope design must take into account three factors.
+ First. we with to maximize observing elliciency Ey.," First, we with to maximize observing efficiency ${\cal E}_S$."
+ Second. we wish to achieve the highest possible signal-to-noise ratio.," Second, we wish to achieve the highest possible signal-to-noise ratio."
+ Third. we must avoil anv distortion problems with the optics.," Third, we must avoid any distortion problems with the optics."
+ There are four effects through which aperture size can impact these factors., There are four effects through which aperture size can impact these factors.
+ Two of these effects. observing ellicieney ancl scintillation noise. will drive us to larger telescopes. while (he other {wo elfects. sky noise and focal plane distortion. will drive us to smaller telescopes.," Two of these effects, observing efficiency and scintillation noise, will drive us to larger telescopes, while the other two effects, sky noise and focal plane distortion, will drive us to smaller telescopes."
+ As we show below. for the observing parameters we have specified an aperture of 5 cm ensures a manageable (uid unique) balance between (he various effects.," As we show below, for the observing parameters we have specified an aperture of 5 cm ensures a manageable (and unique) balance between the various effects."
+ Assuming Nyquist sampling. at most half the light from a point source [alls within one pixel.," Assuming Nyquist sampling, at most half the light from a point source falls within one pixel."
+ We can directly. calculate the ratio of time lost to readout Zi44 to the time spent exposing Lox). where WW. is the well depth of the detector pixels. and Wi = 10?e 1 a fiducial well depth.," We can directly calculate the ratio of time lost to readout $T_{\rm read}$ to the time spent exposing $T_{\rm exp}$, where $W$ is the well depth of the detector pixels, and $W_{0}$ = $10^{5}{\rm e}^{-}$ is a fiducial well depth."
+" Note that the factor 10.I which arises [rom the need to avoid saturation of the brightest stas (where Vi,=Vinay— AV). has been broken up into two terms to permit easy comparisons of equation (14)) with equations (16)) ancl (18)) below."," Note that the factor $10^{-0.4(V_{\rm min}-10)}$, which arises from the need to avoid saturation of the brightest stars (where $V_{\rm min} = V_{\rm max} - \Delta V$ ), has been broken up into two terms to permit easy comparisons of equation \ref{equtexp}) ) with equations \ref{equscint2}) ) and \ref{equskynoise}) ) below."
+ In order to maximize the efficiency. Ey. the fractionof observing time devoted to readout should be minimized. aud therefore. according to equation (14)). so should the aperture size.," In order to maximize the efficiency ${\cal E}_{S}$, the fractionof observing time devoted to readout should be minimized, and therefore, according to equation \ref{equtexp}) ), so should the aperture size."
+ The telescope will operate reasonably efficiently so long as TiaS Tu., The telescope will operate reasonably efficiently so long as $T_{\rm read} \la T_{\rm exp}$ .
+of po.,of $\rho_0$.
+ In comparing mocdel results we therefore note that a change of po may also imply a change in eo., In comparing model results we therefore note that a change of $\rho_0$ may also imply a change in $a_0$.
+ Finally. the slope of the power-law energy distribution of the relativistic electrons strongly inlluences the slope of the observed racio spectrum.," Finally, the slope of the power-law energy distribution of the relativistic electrons strongly influences the slope of the observed radio spectrum."
+ We adjust the value of 6 in the range 2 to 2.5 as discussed in the next sub-section., We adjust the value of $\delta$ in the range 2 to 2.5 as discussed in the next sub-section.
+ For cach individual lobe we use the lobe length and aspect ratio as well as the luminosity densities measured at two frequencies in the fitting process., For each individual lobe we use the lobe length and aspect ratio as well as the luminosity densities measured at two frequencies in the fitting process.
+ Table 4. summarises the model inputs derived from our observations., Table \ref{obpara} summarises the model inputs derived from our observations.
+ To determine the model. parameters po. €. / and. ἐς. we randomly choose a large number of combinations of these four parameters and calculate the prediction of the mocel for the lobe length. L and the luminosity densities £7. at two observing frequencies.," To determine the model parameters $\rho_0$, $Q$, $t$ and $t_{\rm s}$, we randomly choose a large number of combinations of these four parameters and calculate the prediction of the model for the lobe length $L$ and the luminosity densities $P_{\nu}$ at two observing frequencies."
+ A given parameter combination is deemed to be consistent with the observations. if the mocel results are all within of the observed. values of £L and the two values of P.," A given parameter combination is deemed to be consistent with the observations, if the model results are all within of the observed values of $L$ and the two values of $P_{\nu}$ ."
+ Each lobe is fitted individually and the results for the source age / and the jet power Q are shown in the top panel of Fig., Each lobe is fitted individually and the results for the source age $t$ and the jet power $Q$ are shown in the top panel of Fig.
+ 4 where we have set à=2 for all obes., \ref{model} where we have set $\delta =2$ for all lobes.
+ It is re-assuring that for most. possible jet. powers the obe age rellects the lobe size with the outer lobes the oldest and the inner lobes the voungest., It is re-assuring that for most possible jet powers the lobe age reflects the lobe size with the outer lobes the oldest and the inner lobes the youngest.
+ We would expect that the jets inflating cach pair of lobes. outer. middle and. inner. jwe the same jet power and age on both sides of the AGN.," We would expect that the jets inflating each pair of lobes, outer, middle and inner, have the same jet power and age on both sides of the AGN."
+ llence we expect the patches in Fig., Hence we expect the patches in Fig.
+ 4. (top) for a given »ur to show at least some overlap., \ref{model} (top) for a given pair to show at least some overlap.
+ While this is the case or the inner and outer lobe pair. the middle pair shows no overlap at. all.," While this is the case for the inner and outer lobe pair, the middle pair shows no overlap at all."
+ Also. the area in the Q-/ plane allowed. for xh lobes of the outer pair is only small.," Also, the area in the $Q$ $t$ plane allowed for both lobes of the outer pair is only small."
+ This may be a result. of a wrong for slope of the initial power- energy spectrum of choicethe relativisticthe electrons. ὃν for some ofthe lobes.," This may be a result of a wrong choice for the slope of the initial power-law energy spectrum of the relativistic electrons, $\delta$, for some of the lobes."
+ The radio spectra of the middle lobes are steep compared to the spectra of most of the other lobes and may imply a steeper initial energy spectrum., The radio spectra of the middle lobes are steep compared to the spectra of most of the other lobes and may imply a steeper initial energy spectrum.
+ In the interests of exploring the parameter space. in the middle panel of Fig.," In the interests of exploring the parameter space, in the middle panel of Fig."
+ we show the result of changing 6 to 2.5 for the middle lobes., \ref{model} we show the result of changing $\delta$ to 2.5 for the middle lobes.
+ We also change 9 to for the northern. outer lobe and to 2.5 for the southern. 2.2inner lobe.," We also change $\delta$ to 2.2 for the northern, outer lobe and to 2.5 for the southern, inner lobe."
+ Both of these also show at least.slightly steeper radio spectra., Both of these also show at leastslightly steeper radio spectra.
+ This demonstrates the ellect of changing ὁ on the model results., This demonstrates the effect of changing $\delta$ on the model results.
+Ranges and steps for all seven basic parameters of the grid of synthetic spectra are given in. Table 2.,Ranges and steps for all seven basic parameters of the grid of synthetic spectra are given in Table 2.
+ We adopt à common convention of quoting metallicity and enhancement of a—elements in logarithmic units with respect to the solar values., We adopt a common convention of quoting metallicity and enhancement of $\alpha$ --elements in logarithmic units with respect to the solar values.
+ The gravity is 1n logarithmic οσς units., The gravity is in logarithmic cgs units.
+ Details of all calculated parameter combinations are given i Figures 1—3., Details of all calculated parameter combinations are given in Figures 1–3.
+ Spectra are placed in gravity-temperature planes. with metallicity coded by a symbol type.," Spectra are placed in gravity–temperature planes, with metallicity coded by a symbol type."
+ Figure 1 covers the most numerous spectra. tthe ones with no a-enhancement and with microturbulent velocity of 2 km s!.," Figure 1 covers the most numerous spectra, the ones with no $\alpha$ –enhancement and with microturbulent velocity of 2 km $^{-1}$."
+ The computed spectra cover the whole gravity-temperature plane except for hot low-gravity models which are not radiatively stable., The computed spectra cover the whole gravity–temperature plane except for hot low-gravity models which are not radiatively stable.
+ Low-temperature spectra (Tay<5000 K) were computed for a sparser set of metallicities due to large requirements of computing time., Low–temperature spectra $T_\mathrm{eff} < 5000$ K) were computed for a sparser set of metallicities due to large requirements of computing time.
+ These spectra will be added online when completed., These spectra will be added online when completed.
+ Fig., Fig.
+ 2 corresponds to a-enhanced cases and Fig., 2 corresponds to $\alpha$ –enhanced cases and Fig.
+ 3 to those with a different value of microturbulent velocity., 3 to those with a different value of microturbulent velocity.
+ Note that each of the symbols actually corresponds to HL (Fey<7000 K) or 14 (£y77000 K) spectra with different values of rotational velocity (see Table 1) and to three different resolving powers., Note that each of the symbols actually corresponds to 11 $T_\mathrm{eff} < 7000$ K) or 14 $T_\mathrm{eff} \ge 7000$ K) spectra with different values of rotational velocity (see Table 1) and to three different resolving powers.
+ All spectra are available as ascii files grouped into different directories according to their resolving power and temperature., All spectra are available as ascii files grouped into different directories according to their resolving power and temperature.
+ The filenames are in a standard format identified in Table 3., The filenames are in a standard format identified in Table 3.
+ So corresponds to a flux calibrated spectrum between 7650 and 8750 Á.. with Vio=10 km s!. 4/A4=20000. [M/H]=-0.5. 5250 K. logg= 4.5.€= 2km «τὶς and no a—enhancement.," So corresponds to a flux calibrated spectrum between 7650 and 8750 , with $V_\mathrm{rot} = 10 $ km $^{-1}$, $\lambda / \Delta \lambda = 20\,000$, $[\mathrm{M} / \mathrm{H} ] = -0.5$, $T_\mathrm{eff} = 5250$ K, $\log g = 4.5$, $\xi = 2$ km $^{-1}$, and no $\alpha$ –enhancement."
+ The calculated grid is by far too large to present all of its properties here. so we explore only sample cross-sections across the grid.," The calculated grid is by far too large to present all of its properties here, so we explore only sample cross-sections across the grid."
+ Figure 4+ 1s a greyscale presentation of the spectra which were normalized to enhance line visibility., Figure 4 is a greyscale presentation of the spectra which were normalized to enhance line visibility.
+ Each panel shows variation along one parameter axis. starting from a spectrum of a non-rotating KO V type star.," Each panel shows variation along one parameter axis, starting from a spectrum of a non-rotating K0 V type star."
+ Note that all spectra were calculated in a wider wavelength domain. but only the 8400-8750 range is plotted for clarity.," Note that all spectra were calculated in a wider wavelength domain, but only the 8400–8750 range is plotted for clarity."
+ The temperature panel of Fig., The temperature panel of Fig.
+ 4 clearly shows the importance of sharp Ca II lines for any radial velocity study., 4 clearly shows the importance of sharp Ca II lines for any radial velocity study.
+ The panel is a textbook example of the expected behaviour of the Paschen lines and metallic lines., The panel is a textbook example of the expected behaviour of the Paschen lines and metallic lines.
+ The metallicity panel illustrates that the Ca II lines remain strong even at the lowest metallicities and the gravity panel shows their presence in all luminosity classes., The metallicity panel illustrates that the Ca II lines remain strong even at the lowest metallicities and the gravity panel shows their presence in all luminosity classes.
+ The rotational velocity and resolving power panels show how the lines get smeared at high rotational velocities or if observing at low resolving powers., The rotational velocity and resolving power panels show how the lines get smeared at high rotational velocities or if observing at low resolving powers.
+ The steps in the calculated grid are relatively small. but the coverage Is not continuous.," The steps in the calculated grid are relatively small, but the coverage is not continuous."
+ As an example. the step in temperature is 250 K (for Tyx10000 K).," As an example, the step in temperature is 250 K (for $T_\mathrm{eff} \le 10\,000$ K)."
+ This is larger than the baselined accuracy of temperature determination for both GAIA and RAVE surveys., This is larger than the baselined accuracy of temperature determination for both GAIA and RAVE surveys.
+ So the grid will have to be interpolated to smaller steps., So the grid will have to be interpolated to smaller steps.
+ Figure 5 illustrates the errors introduced by a simple linear interpolation., Figure 5 illustrates the errors introduced by a simple linear interpolation.
+ At a certain grid point ¢ with the parameter values p; we compare the true synthetic spectrum S(p;) with the spectrum $ obtained from a linear combination of spectra at neighbouring grid. points: So=fiASQ+SIG)., At a certain grid point $i$ with the parameter values $p_i$ we compare the true synthetic spectrum $S(p_i)$ with the spectrum $S'$ obtained from a linear combination of spectra at neighbouring grid points: $S' = f_{i-1} S(p_{i-1}) + f_{i+1} S(p_{i+1})$.
+" The weights f, and fi, are optimized so that [Spa—§’Fda is minimal.", The weights $f_{i-1}$ and $f_{i+1}$ are optimized so that $\int [S(p_i) - S']^2 d\lambda$ is minimal.
+ The difference between the interpolated values of parameters p’ and the true ones p; can then be expressed in units of a grid step: Figure 5 shows that linear interpolation is accurate to =10 oof the grid step., The difference between the interpolated values of parameters $p'$ and the true ones $p_i$ can then be expressed in units of a grid step: Figure 5 shows that linear interpolation is accurate to $\simlt 10$ of the grid step.
+ Note that this is the worst case scenario. corresponding to a reconstruction of the spectrum at the middle of the grid interval.," Note that this is the worst case scenario, corresponding to a reconstruction of the spectrum at the middle of the grid interval."
+ Linear interpolation would be more accurate for spectra lying closer to one of the grid points., Linear interpolation would be more accurate for spectra lying closer to one of the grid points.
+ The results could be improved further by employing non-linear interpolation schemes., The results could be improved further by employing non-linear interpolation schemes.
+ One may conclude that linear interpolation itself does not introduce errors exceeding 25 K in temperature (for Typ<10000 K). 0.05 dex in [M/H] or logg and | km s! in V4.," One may conclude that linear interpolation itself does not introduce errors exceeding 25 K in temperature (for $T_\mathrm{eff} < 10\,000$ K), 0.05 dex in $[\mathrm{M}/\mathrm{H}]$ or $\log g$ and 1 km $^{-1}$ in $V_\mathrm{rot}$."
+ Note that other errors are more important: degeneracy of parameter values fitting spectra with a limited signal to noise ratio complicates their determination (Bailer-Jones 2003. see also Fig.," Note that other errors are more important: degeneracy of parameter values fitting spectra with a limited signal to noise ratio complicates their determination (Bailer-Jones 2003, see also Fig."
+ | in Zwitter 2002)., 1 in Zwitter 2002).
+ Also. spectra of real starsdo not correspond exactly to the synthetic spectra due to their peculiarities eemission lines. varied abundances of individual elements. non-LTE effects. and atmospheric structure).," Also, spectra of real starsdo not correspond exactly to the synthetic spectra due to their peculiarities emission lines, varied abundances of individual elements, non-LTE effects, and non-static atmospheric structure)."
+"prominent for intermediate harmonics, we nonetheless follow the recommendation of ? to flag this mode as suspect.","prominent for intermediate harmonics, we nonetheless follow the recommendation of \citet{Gilliland} to flag this mode as suspect."
+" The quasi-regularity of the small (6v,9) and large frequency separations (Av,o and Av?) is evident from Figs.", The quasi-regularity of the small $\delta\nu_{n0}$ ) and large frequency separations $\Delta\nu_{n0}$ and $\Delta\nu_{n2}$ ) is evident from Figs.
+ [7] and [8]., \ref{Mulder_echelle} and \ref{Scully_echelle}.
+" Notice that if these stars were to strictly obey the asymptotic relation in Eq. [I],"," Notice that if these stars were to strictly obey the asymptotic relation in Eq. \ref{asymptotic},"
+ they would then exhibit vertical ridges in the écchelle diagram provided use of the correct Av., they would then exhibit vertical ridges in the écchelle diagram provided use of the correct $\Delta\nu$.
+" The small separation óv,o is, however, more clearly distinguished in the case of KIC 10920273, which might be an indication of smaller mode linewidths in this cooler star (see Sect."," The small separation $\delta\nu_{n0}$ is, however, more clearly distinguished in the case of KIC 10920273, which might be an indication of smaller mode linewidths in this cooler star (see Sect."
+ ?? fora discussion on mode linewidths)., \ref{WH} for a discussion on mode linewidths).
+" A striking feature in both écchelle diagrams is the jagged appearance of the /=1 ridge, a trademark of the presence of avoided crossings and an indicator of the evolved nature of these stars."," A striking feature in both écchelle diagrams is the jagged appearance of the $l\!=\!1$ ridge, a trademark of the presence of avoided crossings and an indicator of the evolved nature of these stars."
+" These same features have also been seen in the cases of ground-based observations of 7 Boo (?),, 6 Hyi (?) and possibly Procyon (?), as well as in the cases of the target HD 49385 (?),, and KASC survey targets KIC 11026764 (?),, KIC 11395018 and KIC 11234888 (?).."," These same features have also been seen in the cases of ground-based observations of $\eta$ Boo \citep{Kjeldsen03}, $\beta$ Hyi \citep{Bedding07} and possibly Procyon \citep{Procyon}, as well as in the cases of the target HD 49385 \citep{HD49385}, and KASC survey targets KIC 11026764 \citep{Gemma}, KIC 11395018 and KIC 11234888 \citep{FurryMathur}."
+" Figure D] displays a so-called p-g diagram as introduced by ?,, where the frequencies of the avoided crossings (i.e., the frequencies of the pure g modes in the core cavity) for a number of stars are plotted against the large separation of the p modes."," Figure \ref{pg} displays a so-called p-g diagram as introduced by \citet{Bedding_pg}, where the frequencies of the avoided crossings (i.e., the frequencies of the pure g modes in the core cavity) for a number of stars are plotted against the large separation of the p modes."
+" Much of the diagnostic potential of mixed modes can be captured in this way, since their overall pattern is determined by the mode bumping at each avoided crossing, which in turn is determined by the g modes trapped in the core."," Much of the diagnostic potential of mixed modes can be captured in this way, since their overall pattern is determined by the mode bumping at each avoided crossing, which in turn is determined by the g modes trapped in the core."
+ This diagram could prove to be an instructive way to display results of many stars and to allow for a first comparison with theoretical models., This diagram could prove to be an instructive way to display results of many stars and to allow for a first comparison with theoretical models.
+" We also report here the possible presence of a/=2 mixed mode in the power spectrum of KIC 10920273 (at 873.10 μΗΖ) that should, however, be confirmed by stellar models."," We also report here the possible presence of a $l\!=\!2$ mixed mode in the power spectrum of KIC 10920273 (at $873.10\:{\rm{\mu Hz}}$ ) that should, however, be confirmed by stellar models."
+ Detection of /=3 modes with photometric observations is made very difficult due to geometric cancellation effects., Detection of $l\!=\!3$ modes with photometric observations is made very difficult due to geometric cancellation effects.
+ Solar-like oscillations with /23 from photometry have nonetheless been reported for a set of low-luminosity red giants by ?.., Solar-like oscillations with $l\!=\!3$ from photometry have nonetheless been reported for a set of low-luminosity red giants by \citet{Bedding_rg}.
+ ? have also reported the presence of /23 modes for the target HD 49385., \citet{HD49385} have also reported the presence of $l\!=\!3$ modes for the target HD 49385.
+" We should bear in mind that, except for ORK and SYD, all the remaining fitters used deterministicmodels in their frequency-domain representations of the data that only contained modes of degree up to /=2, meaning that a statistical assessment of the presence or not of /=3 modes could not be done."," We should bear in mind that, except for ORK and SYD, all the remaining fitters used deterministicmodels in their frequency-domain representations of the data that only contained modes of degree up to $l\!=\!2$, meaning that a statistical assessment of the presence or not of $l\!=\!3$ modes could not be done."
+" ORK and SYD, which were the only fitters that did not make any prior assumptions about the degree of the modes, have not reported the detection of modes that could be interpreted as /=3 modes."," ORK and SYD, which were the only fitters that did not make any prior assumptions about the degree of the modes, have not reported the detection of modes that could be interpreted as $l\!=\!3$ modes."
+" A thorough discussion of the mode linewidths, heights, and amplitudes goes beyond the scope of this work."," A thorough discussion of the mode linewidths, heights, and amplitudes goes beyond the scope of this work."
+" However, there are some aspects we would like to mention here."," However, there are some aspects we would like to mention here."
+ The intrinsic frequency resolution of the spectra (x0.05 uHz) makes it possible to resolve the modes., The intrinsic frequency resolution of the spectra $\approx\!0.05\:{\rm{\mu Hz}}$ ) makes it possible to resolve the modes.
+ This condition is obeyed provided the observation length T>2Tmode (?).., This condition is obeyed provided the observation length $T\!\gg\!2\tau_{\rm{mode}}$ \citep{Chaplin03}.
+" Figure displays, for each star, the linewidths of the radial modes ποreturned by the respective (see also Tables and [8))."," Figure \ref{widths} displays, for each star, the linewidths of the radial modes returned by the respective (see also Tables \ref{Freq_Mulder2} and \ref{Freq_Scully2}) )."
+ The radial modes considered are those belonging to the set.., The radial modes considered are those belonging to the .
+" Notice the near-constancy with frequency of the mode linewidths in the case of KIC 10920273, whereas for KIC 10273246 the linewidths increase steadily,"," Notice the near-constancy with frequency of the mode linewidths in the case of KIC 10920273, whereas for KIC 10273246 the linewidths increase steadily,"
+Before producing our final cluster catalog. we need to define a criferiun to identify1 clusters by associating Huctuations detected in different inagnitude bius.,"Before producing our final cluster catalog, we need to define a criterium to identify clusters by associating fluctuations detected in different magnitude bins."
+ The criteri will consists of à maxima projected distauce vetween the centers of fluctuations to be associated aud of a minimi uunber of coincident fluctuations required or a positive detection. [Ng," The criterium will consists of a maximum projected distance between the centers of fluctuations to be associated and of a minimum number of coincident fluctuations required for a positive detection, $N_{min}$."
+ In fact. because of the statistical noise of the oreeround/backeround galaxy distribution. the ceuters of the 8ctuations produced by a cluster in different magnitude bins will be sliehtlv different.," In fact, because of the statistical noise of the foreground/background galaxy distribution, the centers of the fluctuations produced by a cluster in different magnitude bins will be slightly different."
+ As far as the umber of coincident fluctuations produced by a cluster is concerned. it will depeud on the cluster distance. richness and luuinosity fiction.," As far as the number of coincident fluctuations produced by a cluster is concerned, it will depend on the cluster distance, richness and luminosity function."
+ Clearly the choice of the criterium has to be mace having in wind the goal of the detection algoritlin and the characteristics of the galaxy catalog., Clearly the choice of the criterium has to be made having in mind the goal of the detection algorithm and the characteristics of the galaxy catalog.
+ Tn order to define the criterium of association for our test application. we peform extended tests both ou he PDCS field aud ou simulations of Poissonian fields with embedded: simulated clusters.," In order to define the criterium of association for our test application, we peform extended tests both on the PDCS field and on simulations of Poissonian fields with embedded simulated clusters."
+ Our simulated fields lave the sale eeneral properties of the PDCS ποια., Our simulated fields have the same general properties of the PDCS field.
+ Tn particular. we run the VGCF on LOO catalogs cach containing LS simulated clusters (a umuboer similar to the iunber of clusters in the PDCS field) eiibedded within a Poissonian galaxy field.," In particular, we run the VGCF on 100 catalogs each containing 18 simulated clusters (a number similar to the number of clusters in the PDCS field) embedded within a Poissonian galaxy field."
+ Clusters are smniulate« as in the xevious subsection., Clusters are simulated as in the previous subsection.
+ All simmlated catalogs coutaiu 25132 ealaxies. ie. the same number of PDCS galaxies.," All simulated catalogs contain 25432 galaxies, i.e. the same number of PDCS galaxies."
+ As a result of our tests. we conskler coincident wo fluctuations with centers separated on the skv by a projected. distance dyxODRy. Ro). where Ry and Ro are the radi of the two fluctuations.," As a result of our tests, we consider coincident two fluctuations with centers separated on the sky by a projected distance $d_{12} \leq
+0.3\,min(R_1,R_2)$ , where $R_1$ and $R_2$ are the radii of the two fluctuations."
+ A ighter criteri would break the sequence of fluctuations corresponding to a real cluster. a looser criterimu would incorrectly associate fluctuations produced by adjacenut clusters/fluctuations to the same cluster.," A tighter criterium would break the sequence of fluctuations corresponding to a real cluster, a looser criterium would incorrectly associate fluctuations produced by adjacent clusters/fluctuations to the same cluster."
+ We now set the nünmuuu nmuuber of fluctuations. Minin required. for the detection. of a cluster.," We now set the minimum number of fluctuations, $N_{min}$, required for the detection of a cluster."
+ In Fig., In Fig.
+ G we plot the average nuuniber of spurious fluctuations ni-ideuti&ed as clusters as a function of μι, 6 we plot the average number of spurious fluctuations mis-identified as clusters as a function of $N_{min}$.
+ The uuuber of spurious clusters drops dramatically as κρεμ increases from 1 to 5., The number of spurious clusters drops dramatically as $N_{min}$ increases from 1 to 5.
+" For N,,;,—5 here are ou average 1.5 spurious clusters per field.", For $N_{min} = 5$ there are on average 1.5 spurious clusters per field.
+ This απ. decreases slowly as ορ ducreases further., This number decreases slowly as $N_{min}$ increases further.
+" This result indicates that N,,;,=5 is aa good choice to keep the number of Poisson fluctuations low while still being scusitive to poor or distant clusters."," This result indicates that $N_{min} =
+5$ is a a good choice to keep the number of Poisson fluctuations low while still being sensitive to poor or distant clusters."
+ In Fig., In Fig.
+" 7 we show the variation with N,,;, of the detection efficiency of simulated clusters at redshifts z=0.3 (panel a). 0.5 (pancl b). aud 0.8 (panel ο)."," 7 we show the variation with $N_{min}$ of the detection efficiency of simulated clusters at redshifts z=0.3 (panel ), 0.5 (panel ), and 0.8 (panel )."
+ At each redshift the different curves correspond to riclinesses ranging from Np=10 to Ngσυ., At each redshift the different curves correspond to richnesses ranging from $N_R=10$ to $N_{R}=60$.
+ Clearly the fraction of detected clusters decreases as Vy) Increases., Clearly the fraction of detected clusters decreases as $N_{min}$ increases.
+ We note that the curve in Fig., We note that the curve in Fig.
+ 6 corresponds to au evaluation of the False Positive Rate (FPR). ic. the probability of detecting as a cluster a random fluctuation of the galaxy distribution.," 6 corresponds to an evaluation of the False Positive Rate (FPR), i.e. the probability of detecting as a cluster a random fluctuation of the galaxy distribution."
+" For [N,,;,,=5. our FPR is very low."," For $N_{min}=5$, our FPR is very low."
+ By comparison. P96 eive. 1.2 spurious detections per square deerce with peak signal ereatcr than 3o.," By comparison, P96 give 4.2 spurious detections per square degree with peak signal greater than $3\sigma$."
+ The curves in Fig., The curves in Fig.
+ 7 correspond to a measure of the False Negative Rate (CENR)., 7 correspond to a measure of the False Negative Rate (FNR).
+" We present here the results of the run of the VOCE ou he V, catalog of the PDCS field at a=13 26"" and à = 20"" 52 (12000).", We present here the results of the run of the VGCF on the $V_4$ catalog of the PDCS field at $\alpha = 13^{h}$ $^{m}$ and $\delta$ = $29^{o}$ 52' (J2000).
+ As diseussed in the previous section. we run the VOCE in bins two magnitude wide. “sliding” with Kl magnitude steps within the magnitude rauge 18.00<VQox23.8.," As discussed in the previous section, we run the VGCF in bins two magnitude wide, “sliding” with 0.1 magnitude steps within the magnitude range $18.00 \leq V_{4} < 23.8$ ."
+ In total we run the VOCF in 39 magnitude ius ancl identify as clusters at least five fluctuations that. according to our criteria. are anenlarly coicideut (sec xevious subsection).," In total we run the VGCF in 39 magnitude bins and identify as clusters at least five fluctuations that, according to our criterium, are angularly coincident (see previous subsection)."
+ Ou a Sun ULTRASpare 30 workstation. the time required to run the whole procedure is about LO minutes.," On a Sun ULTRASparc 30 workstation, the time required to run the whole procedure is about 40 minutes."
+ Our output cluster list consists of 37 objects., Our output cluster list consists of 37 objects.
+ We characterize each cluster with the properties of the fluctuation with the lighest signal-to-noise ratio as estiniatec w the ratio of the numuber of cluster galaxies to the square root of the number of background galaxies expected within the cluster area., We characterize each cluster with the properties of the fluctuation with the highest signal-to-noise ratio as estimated by the ratio of the number of cluster galaxies to the square root of the number of background galaxies expected within the cluster area.
+ Iu Table 1. for each cluster we list: 1) ideutification munhber. 2) J2000 right ascension aud 3) J2000 declination. 1) radius. 5) the cluster signal-to-noise ratio. 6) the estimated umber of cluster ealaxies aud 7) the uuuber of expected backeromn ealaxies. 8) the central maguitude of the bin where we detect the cluster with the highest sigual-to-nolse ratio. 9) he cross-identification with the PDCS catalog.," In Table 1, for each cluster we list: 1) identification number, 2) J2000 right ascension and 3) J2000 declination, 4) radius, 5) the cluster signal-to-noise ratio, 6) the estimated number of cluster galaxies and 7) the number of expected background galaxies, 8) the central magnitude of the bin where we detect the cluster with the highest signal-to-noise ratio, 9) the cross-identification with the PDCS catalog."
+ Iu Fig., In Fig.
+ δ we plot circles on the sky corresponding to our clusters (solid lines)., 8 we plot circles on the sky corresponding to our clusters (solid lines).
+" We label oux clusters with a ""V7 followed by their order umuuber in Table 1.", We label our clusters with a “V” followed by their order number in Table 1.
+ In Fie., In Fig.
+ 9. we eive a evaphic sumuuary of all the fluctuations of cach cluster.," 9, we give a graphic summary of all the fluctuations of each cluster."
+ The abscissa is the order umber of the uaenitude bin of the fluctuation and each cluster is represeuted. by a row of circles parallel to the magnitude bin axis., The abscissa is the order number of the magnitude bin of the fluctuation and each cluster is represented by a row of circles parallel to the magnitude bin axis.
+ The radii of the circles are scaled with the signal-to-noise ratio of the detection., The radii of the circles are scaled with the signal-to-noise ratio of the detection.
+ We renuud rere that we fit a circle to a fluctuation after its detection. aud that the oulv use of the circle is to provide a convenieut wav to catalog the cluster with a ceuter and a radius.," We remind here that we fit a circle to a fluctuation after its detection, and that the only use of the circle is to provide a convenient way to catalog the cluster with a center and a radius."
+ The richest aud more reliable clusters exhibit. in Fie.," The richest and more reliable clusters exhibit, in Fig."
+ 9. a sequence of fluctuations.," 9, a sequence of fluctuations."
+ The signal-to-noise ratio of these fluctuations regularly increases up to a maxiumni and then decreases., The signal-to-noise ratio of these fluctuations regularly increases up to a maximum and then decreases.
+ Several fainter clusters show the same behavior. although at a generally lower S/N level.," Several fainter clusters show the same behavior, although at a generally lower S/N level."
+ Clearly. the position of he fluctuations along the magnitude axisis related to the cluster distance.," Clearly, the position of the fluctuations along the magnitude axisis related to the cluster distance."
+ Some clusters. for example VF aud V12. displav substantial gaps along the sequence.," Some clusters, for example V7 and V12, display substantial gaps along the sequence."
+ The suspicion is, The suspicion is
+he remnant lifetime eenerally iucreases with the density and the flow velocity of the surrouncdiug nedium.,the remnant lifetime generally increases with the density and the flow velocity of the surrounding medium.
+ The üighest value found iu our simulations. Mf=77Mvr. is obtained when a=10°cu Sand e=5000kims!.," The highest value found in our simulations, $Mt = 77 \Msol \yr$, is obtained when $n = 10^{6} \pcm3$ and $v = 5000 \kmps$."
+ A supernova rate of L/vr would then imply a mass or the clouds enüttiug the IIILs of up to ~soXl... This is easily compatible with the lower cud of DELR πάσα estimates in the literature(c.e... Peterson 1997)). althoueh our model (and most others} would ο severely challenged to explain uch more extreme estimates of the uass of BELR eas (see Balelwin 2003. aac references herein).," A supernova rate of 1/yr would then imply a mass for the clouds emitting the HILs of up to $\sim 80 \Msol$, This is easily compatible with the lower end of BELR mass estimates in the literature, Peterson \cite{P1997}) ), although our model (and most others) would be severely challenged to explain much more extreme estimates of the mass of BELR gas (see Baldwin \cite{B2003} and references therein)."
+" We note that it is currently uuclear row this mass 1s partitioned between the TIL aud LIL eas in these higher estimates,", We note that it is currently unclear how this mass is partitioned between the HIL and LIL gas in these higher estimates.
+ Iu earlier work it was shown that for typical QSO paraucters the power going into supernova renimauts is comparable to that of the QSO wind. but is 11ncli less than he bolometric QSO luminosity (Perry Dyson 1985).," In earlier work it was shown that for typical QSO parameters the power going into supernova remnants is comparable to that of the QSO wind, but is much less than the bolometric QSO luminosity (Perry Dyson \cite{PD1985}) )."
+ However. it is more difficult to estimate whether emission roni the SN would be visible above the QSO in a specific wavebaud.," However, it is more difficult to estimate whether emission from the SN would be visible above the QSO in a specific waveband."
+ The typical J-baud magnitude for a QSO at a redshift +~ lods c91δ19% whereas the J1II baud naenitude for a type Ia SN is ~21 at comparable :.," The typical J-band magnitude for a QSO at a redshift $z \sim 1$ is $\sim 18-19$, whereas the J+H band magnitude for a type Ia SN is $\sim 24$ at comparable $z$."
+ Ou his basis. individual SN will no be disceruible. but clearly his conclusion depends on the luminosity of the QSO. as well as other variables such as the oricutation of the SN with respect o the molecular torus. and the ambieut density of the surroundings(e.g... if the SNR expands iuto a nearby noleculu cloud then its luninosity could be senuificautlv increased).," On this basis, individual SN will not be discernible, but clearly this conclusion depends on the luminosity of the QSO, as well as other variables such as the orientation of the SN with respect to the molecular torus, and the ambient density of the surroundings, if the SNR expands into a nearby molecular cloud then its luminosity could be significantly increased)."
+ Detailed uunuerical modelling will be required to determine the likelihood of this possibility., Detailed numerical modelling will be required to determine the likelihood of this possibility.
+ One of the most interesting questions concerning ACGNs is the connection between nuclear aud starburst activity., One of the most interesting questions concerning AGNs is the connection between nuclear and starburst activity.
+ Miwed stayburst-ACUN. SOMYCOS ay c, Mixed starburst-AGN sources may be
+and with indices a=0.2*27 and 8=1.4707.,and with indices $\alpha = 0.2^{+0.2}_{-0.1}$ and $\beta = 1.4^{+0.3}_{-0.6}$.
+" The broken power law fit is very good, giving X?=0.63 for 4 d.o.f."," The broken power law fit is very good, giving $\chi^2 = 0.63$ for 4 d.o.f."
+ This result agrees with previous studies (e.g.Daigneetal.Piran 2007)..," This result agrees with previous studies \citep[e.g.][]{Daigne(2006),Guetta(2007)}."
+" The cannot be fitted with a single power law, since such fit give high y?, rejecting such a model with high significance (98%))."," The cannot be fitted with a single power law, since such fit give high $\chi^2$, rejecting such a model with high significance )."
+ This contradict the results of Pélangeonetal.(2008) who studied the HETE-2 GRBs and found a consistency with a single power lawfunction., This contradict the results of \cite{Pelangeon(2008)} who studied the HETE-2 GRBs and found a consistency with a single power law.
+". The rate is described as well by a broken power law for 14- z, with a break at 2=3.1702 and indices of ny=2.1702 and πο=—1.4771."," The rate is described as well by a broken power law for $1+z$ , with a break at $z = 3.1^{+0.6}_{-0.8}$ and indices of $n_1 = 2.1^{+0.5}_{-0.6}$ and $n_2 = -1.4^{+2.4}_{-1.0}$."
+" The local event rate is ppc1.3*52[Gpc ?yr-!], in agreement with previous studies e.g. Schmidt (1999)[[po~ 1.5] (see however Schmidt (2001b)[[po~ 0.15)]); Guettaetal.(2005)[[po 0.5]; GuettaandDellaValle (2007)[[po 1.1]; Liangetal. (2007)[[oo~ 1.1]; al. 2008)[[po.Z 0.5]."," The local event rate is $\rho_0 \simeq 1.3^{+0.6}_{-0.7} [Gpc^{-3}yr^{-1}]$ , in agreement with previous studies e.g. \cite{Schmidt(1999)}[ $\rho_0 \simeq 1.5$ ] (see however \cite{Schmidt(2001b)}[ $\rho_0 \simeq 0.15$ )]); \cite{Guetta(2005)}[ $\rho_0 \simeq 0.5$ ]; \cite{GuettaDV(2007)}[ $\rho_0 \simeq 1.1$ ]; \cite{Liang(2007)}[ $\rho_0 \simeq 1.1$ ]; \cite{Pelangeon(2008)}[ $\rho_0 \gtrsim 0.5$ ]."
+" The main factors determining thelow redshift (current) event rateare the overall GRB rate normalization, the low-redshift slope, the low end of the"," The main factors determining thelow redshift (current) event rateare the overall GRB rate normalization, the low-redshift slope, the low end of the"
+tomogram is then at the right phase to correspond. to emission from either the secondary star or the bright spot where the stream hits the cise (or a mixture of these: the hase uncertainty prohibits a secure distinction between the »ossibilities).,tomogram is then at the right phase to correspond to emission from either the secondary star or the bright spot where the stream hits the disc (or a mixture of these; the phase uncertainty prohibits a secure distinction between the possibilities).
+ Weaker emission is then seen looping leftsvarcds owards the higher-velocity feature. and. could be emission rom the overllowing stream.," Weaker emission is then seen looping leftwards towards the higher-velocity feature, and could be emission from the overflowing stream."
+ We don't have sullicient information (masses and inclination) to interpret the velocities in. the tomogram directly. but we can perform a plausibility check.," We don't have sufficient information (masses and inclination) to interpret the velocities in the tomogram directly, but we can perform a plausibility check."
+ The high-velocity. wings in eclipsing SW Sex stars extend. to velocities of 14001600 iin stars such as SW Sex itself (Dhillon. Marsh JJones 1997) and V1315 λα (Llellier 1996): the equivalent component in eextends to 1300 in data with a comparable signal-to-noise ratio.," The high-velocity wings in eclipsing SW Sex stars extend to velocities of 1400–1600 in stars such as SW Sex itself (Dhillon, Marsh Jones 1997) and V1315 Aql (Hellier 1996); the equivalent component in extends to 1300 in data with a comparable signal-to-noise ratio."
+ These velocities match if hhas an inclination of60.. or a sin/ of O.S7.," These velocities match if has an inclination of, or a $\sin i$ of 0.87."
+ Further. adopting à white-dwarl mass of 0.7 aancd a red-cdwarl mass of 0.1 iimplies that the red. cwarf has an orbital velocity of 4405. that the Lagrangian point orbits at 290|. and that the outer edge of the dise (assuming it is located at the tidal limit) orbits at. 650 (μου Warner 1995. chapter 2. for the relevant formulae).," Further, adopting a white-dwarf mass of 0.7 and a red-dwarf mass of 0.1 implies that the red dwarf has an orbital velocity of 440, that the Lagrangian point orbits at 290, and that the outer edge of the disc (assuming it is located at the tidal limit) orbits at 650 (see Warner 1995, chapter 2, for the relevant formulae)."
+ These values compare with the observed S-wave amplitude of 350L|. or 400 wwith the above sini.," These values compare with the observed S-wave amplitude of 350, or 400 with the above $\sin i$."
+ Thus the S-wave is compatible with arising [rom the secondary or the carly part of the stream. but less compatible with arising [rom the streamdisc impact (unless the inclination or the white-clwarl mass are lower than adopted. above).," Thus the S-wave is compatible with arising from the secondary or the early part of the stream, but less compatible with arising from the stream–disc impact (unless the inclination or the white-dwarf mass are lower than adopted above)."
+ Thus. overall. the line profiles are consistent with the stream-overllow idea. in that both the lower-velocity S-wave and the lino wings have compatible velocities.," Thus, overall, the line profiles are consistent with the stream-overflow idea, in that both the lower-velocity S-wave and the line wings have compatible velocities."
+ A crucial observation for models of iis that the A-ray lighteurve. varies only with the 2147-5 spin period. and not with the orbital evcle nor the orbital sidebancs of the spin period. (Llellier 11998).," A crucial observation for models of is that the X-ray lightcurve varies only with the 2147-s spin period, and not with the orbital cycle nor the orbital sidebands of the spin period (Hellier 1998)."
+ This implies that the accreting material loses knowlecdee of orbital phase before attaching to field. lines., This implies that the accreting material loses knowledge of orbital phase before attaching to field lines.
+ ‘This. in turn. suggests that. if the stream-overllow mocel is correct. the overflowing stream does not travel far enough o encounter the magnetosphere. but. instead. re-inipacts he disc further out.," This, in turn, suggests that, if the stream-overflow model is correct, the overflowing stream does not travel far enough to encounter the magnetosphere, but instead re-impacts the disc further out."
+ This contrasts with suggestions [or other IPs. for instance FO Aqr. where the interaction of he overllowing stream with the magnetopshere was invoked specifically to explain an X-ray beat pulse (οπου 1993: Bearcmore 1998).," This contrasts with suggestions for other IPs, for instance FO Aqr, where the interaction of the overflowing stream with the magnetopshere was invoked specifically to explain an X-ray beat pulse (Hellier 1993; Beardmore 1998)."
+ Two caveats should. be mace., Two caveats should be made.
+ First. the overllow müsght be intermittent. and might not have been occurring during the X-ray observations.," First, the overflow might be intermittent, and might not have been occurring during the X-ray observations."
+ Indeed. the X-ray beat pulse in FO Aqr is variable and sometimes absent.," Indeed, the X-ray beat pulse in FO Aqr is variable and sometimes absent."
+ Secondly. we should consider whether the 2147-5 period is misidentified. ancl is actually the beat Q)) period.," Secondly, we should consider whether the 2147-s period is misidentified, and is actually the beat ) period."
+ This. though. would imply a spin period of 1509 s. and no such periodicity has ever been seen inCen: and further. other observed »riodicities. such as the IS60-5. modulation. would. then lave no natural identification.," This, though, would imply a spin period of 1509 s, and no such periodicity has ever been seen in; and further, other observed periodicities, such as the 1860-s modulation, would then have no natural identification."
+ ‘There is. however. a significant beat-cvele modulation in he line V/1t ratios (Section 2).," There is, however, a significant beat-cycle modulation in the line V/R ratios (Section 2)."
+ Phis can be explained in the standard way for optical beat periods. namely irradiation of structure fixed in the binary [rame (secondary or streani) w the spin-pulsed X-ray beam.," This can be explained in the standard way for optical beat periods, namely irradiation of structure fixed in the binary frame (secondary or stream) by the spin-pulsed X-ray beam."
+ However. the period of the observed: modulation. differs from. the expected value. by LT per cent.," However, the period of the observed modulation differs from the expected value by 0.7 per cent."
+ Over the 2-d span of the observations this amounts to a shift of 0.85 evcles., Over the 2-d span of the observations this amounts to a shift of 0.35 cycles.
+ | possible explanation is that over the 2-cl interval the N-rav. beam switched from illuminating (preclominanth:) the secondary. to illuminating (predominantly) the structure formed where the overflowing stream re-impacts the disc.," A possible explanation is that over the 2-d interval the X-ray beam switched from illuminating (predominantly) the secondary, to illuminating (predominantly) the structure formed where the overflowing stream re-impacts the disc."
+ As can be seen from the tomogram. these two regions are separated by 0.35 in orbital phase.," As can be seen from the tomogram, these two regions are separated by 0.35 in orbital phase."
+ Hf correct. this again suggeests that the overllow is intermittent. occuring only some of the timo.," If correct, this again sugggests that the overflow is intermittent, occuring only some of the time."
+ One puzzle for the above model is the observation of the 1860-5. photometric modulation. identified with O0)) whenw lis not seen.," One puzzle for the above model is the observation of the 1860-s photometric modulation, identified with ), when is not seen."
+ Reprocessing of N-ravs would. likely result. in an optical moclulation. as is observed in many LPs. but not 3).," Reprocessing of X-rays would likely result in an optical modulation, as is observed in many IPs, but not )."
+ One plausible explanation. that the illuminating X-ray beam is double-pealked. resulting in reprocessing at Q)). is contradicted by the fact that the observed. X-ray pulse is nearly sinusoical (Llellier 11998).," One plausible explanation, that the illuminating X-ray beam is double-peaked, resulting in reprocessing at ), is contradicted by the fact that the observed X-ray pulse is nearly sinusoidal (Hellier 1998)."
+ hus. this explanation only works if the X-ray pulse is beamed such that it is double-peaked in the orbital plane but sinusoidal from our line of sight. which is unlikely.," Thus, this explanation only works if the X-ray pulse is beamed such that it is double-peaked in the orbital plane but sinusoidal from our line of sight, which is unlikely."
+ Lavine considered a mocel for bbased on the conventional partial disc. we now consider the alternative discless model based on a diamagnetie Dow.," Having considered a model for based on the conventional partial disc, we now consider the alternative discless model based on a diamagnetic flow."
+ This model was proposed by Ixing (1993) and Wynn Whine (1995). see also Wynn (2001).," This model was proposed by King (1993) and Wynn King (1995), see also Wynn (2001)."
+ Ht treats the accretion How as à set of diamagnetie blobs. and represents these by the particles in a hyvdrodynamical code. with the addition of a magnetic drag term which acts like the tension of the magnetic field lines.," It treats the accretion flow as a set of diamagnetic blobs, and represents these by the particles in a hydrodynamical code, with the addition of a magnetic drag term which acts like the tension of the magnetic field lines."
+ Εις term is proportional to the rate at which particles cross field lines. giving an acceleration where ο and v; are the velocities of the material and field and the symbol L indicates the component perpendicular to the field lines.," This term is proportional to the rate at which particles cross field lines, giving an acceleration where $\mbold{v}$ and $\mbold{v}_{\rm f}$ are the velocities of the material and field and the symbol $\perp$ indicates the component perpendicular to the field lines."
+ Phe parameter & is dependent on factors such as the the local field strength. blob density and Alfvénn," The parameter $k$ is dependent on factors such as the the local field strength, blob density and Alfvénn"
+the times of maxima of the 11-21 variation to obtain au ephemeris wherein both linear aud quadratic terms are sienificaut.,the times of maxima of the 11-m variation to obtain an ephemeris wherein both linear and quadratic terms are significant.
+ To our knowledge. this is the most precise ephemeris in the literature at this time.," To our knowledge, this is the most precise ephemeris in the literature at this time."
+ We detect the periodicity in the ASN liebt curve with Ligh significance: see Fig., We detect the periodicity in the ASM light curve with high significance; see Fig.
+ 12. and. Tabο 3., \ref{fig:pds1820} and Table \ref{tbl:detect}.
+" The period we obtain. P=685.01192+O.00007 s. is an average over the curatiui of the ASAT [250.0005data set aud1] effectively a»plies near the miud-time of this interval. Ίνοι, near MJD 52620."," The period we obtain, $P = 685.01092 \pm 0.00007\,[\pm 0.00054]$ s, is an average over the duration of the ASM data set and effectively applies near the mid-time of this interval, i.e., near MJD 52620."
+ We Call COpare our])oriod iueasureimentwiti the period of D—685.01126 s predicted for this epoch |w the Chou&(αποΊαν(2001), We can compare our period measurementwith the period of $P = 685.01126$ s predicted for this epoch by the \citet{chougrn01} ephemeris.
+ οΣπ. iat the s1ualler ο the wo uncertaiuties in the ASM perio that are given in Table 30 is applicable. then the ASA period is about | standard deviations below the prediction of the Chou&Caindlay(2001). ephemeris.," If we assume that the smaller of the two uncertainties in the ASM period that are given in Table \ref{tbl:detect} is applicable, then the ASM period is about 4 standard deviations below the prediction of the \citet{chougrn01}
+ ephemeris."
+ This sugeests that the cocfiicicut of the quadratic term in that ephemeris should be more negative., This suggests that the coefficient of the quadratic term in that ephemeris should be more negative.
+ However. we have not undertaken a proper joint analvsis of all of the timing results and so this conclusion iust be regarded as tentative.," However, we have not undertaken a proper joint analysis of all of the timing results and so this conclusion must be regarded as tentative."
+ The orbital periods of a umber of IININDs are listed iu Table 3.., The orbital periods of a number of HMXBs are listed in Table \ref{tbl:detect}.
+ Iu contrast with the results on LAINBs. Table 3 shows that the majority of the most significant detections of these svstenis were made in the 5-12 keV energy baud.," In contrast with the results on LMXBs, Table \ref{tbl:detect} shows that the majority of the most significant detections of these systems were made in the 5-12 keV energy band."
+ This is in accordance with expectations based on the observation that πας of the accreting compact objects in IIMXNDs are X-ray pulsars that teud to have relatively hard X-ray spectra., This is in accordance with expectations based on the observation that many of the accreting compact objects in HMXBs are X-ray pulsars that tend to have relatively hard X-ray spectra.
+0535426 is a De/N-rav pulsar that can be rather bright in N-ravs duriug trausicut outbursts., is a Be/X-ray pulsar that can be rather bright in X-rays during transient outbursts.
+ The carly observations are reviewed by Fingerctal.(1996)., The early observations are reviewed by \citet{fwh96}.
+. Pulse timing analyses of observations obtained with the BATSE instrument were used to determine that the orbital period is P?=110.3250.3 davs (Fineeretal.1996.andreferences therein). , Pulse timing analyses of observations obtained with the BATSE instrument were used to determine that the orbital period is $P = 110.3 \pm 0.3$ days \citep[and references therein]{fwh96}. .
+A portion of an ASAI power density spectrun is shown iu Figure 1L., A portion of an ASM power density spectrum is shown in Figure \ref{fig:pdsj1008}. .
+ This figure, This figure
+photoevaporated by the ionising background. and/or ejected from haloes due to supernovae feedback.,"photoevaporated by the ionising background, and/or ejected from haloes due to supernovae feedback."
+" At 2=3. the mass-scale of 107LO""fINS corresponds to a circular velocity of 1015kms'. or a virial temperat slightly below 10!K."," At $z=3$, the mass-scale of $10^8 - 10^{8.5}\himsun$ corresponds to a circular velocity of $10-15\kms$, or a virial temperature slightly below $10^4\,{\rm K}$."
+ Note that this is a smaller mass-scale ure.than was suggested by Quinn.Ixatz.&Efstathiou(1996). and Thoul&Weinberg (1996).. who argued that haloes with circular velocities less than 40kms+ are unlikely to harbour DLAs.," Note that this is a smaller mass-scale than was suggested by \citet{Qui96} and \citet{Tho96}, who argued that haloes with circular velocities less than $40\kms$ are unlikely to harbour DLAs."
+ We will discuss this point further in Section 7.., We will discuss this point further in Section \ref{section:discussion}.
+ 1n Figure 4.. we show DLA cross-sections as a function of total halo mass for >=1 and z— 0. and the parameters of the fitted. power-laws are summarised in Table 4...," In Figure \ref{area_lowz.eps}, we show DLA cross-sections as a function of total halo mass for $z=1$ and $z=0$ , and the parameters of the fitted power-laws are summarised in Table \ref{table:lowzfit}."
+ A similar trend in the slope as a function of resolution exists at 2=1 as we saw ab 2=3., A similar trend in the slope as a function of resolution exists at $z=1$ as we saw at $z=3$.
+ Le is clearthat the slope cannot. be determined. reliably for G4 and G5 at 2=0 (and possibly at 2= 1) clue to limited resolution. as is evident [rom the ‘stripes’ seen at low cross-sections in the bottom two panels of Figure 4..," It is clearthat the slope cannot be determined reliably for G4 and G5 at $z=0$ (and possibly at $z=1$ ) due to limited resolution, as is evident from the `stripes' seen at low cross-sections in the bottom two panels of Figure \ref{area_lowz.eps}."
+ Dark matter haloes with masses below the resolution limit ofa simulation cannot be resolved., Dark matter haloes with masses below the resolution limit of a simulation cannot be resolved.
+ This is a serious problem. when one tries to compute the number density of DLAs based on a cosmological simulation that does not resolve all small mass haloes that may host a DLA., This is a serious problem when one tries to compute the number density of DLAs based on a cosmological simulation that does not resolve all small mass haloes that may host a DLA.
+ Note in particular that the number density of dark matter haloes is known to increase strongly towards lower masses., Note in particular that the number density of dark matter haloes is known to increase strongly towards lower masses.
+ Even a small incompleteness at low masses will hence prevent a reliable estimate of the DLA abundance if only a simple number count of DLAs found in a cosmological simulation is used., Even a small incompleteness at low masses will hence prevent a reliable estimate of the DLA abundance if only a simple number count of DLAs found in a cosmological simulation is used.
+ Vo overcome this limitation. Gardneretal.(1997a.b.2001) convolved a theoretical Gt to the dark matter halo mass function with the measured relationship between DLA cross-section and halo mass.," To overcome this limitation, \citet{Gar97a,Gar97b,Gar01}
+ convolved a theoretical fit to the dark matter halo mass function with the measured relationship between DLA cross-section and halo mass."
+ In this way. they were able to correct for incompleteness in the resolved halo abundance of the simulations.," In this way, they were able to correct for incompleteness in the resolved halo abundance of the simulations."
+ The cumulative abundance (or equivalently the rate of incidence) of DLAs per unit redshift as a function of halo mass in this approach can be expressed as πο. -- ο eppattMA where aayCM.2) ix the dark matter halo. mass function (for whieh we usethe Sheth&Lormen(1999) porameterisation). and dr/dz=ο(ς) with H(z)=οί)=HyQuL2)|OY for a flat universe.," The cumulative abundance (or equivalently the rate of incidence) of DLAs per unit redshift as a function of halo mass in this approach can be expressed as (>M, z) = (M',z) M', where $n_{\rm dm}(M,z)$ is the dark matter halo mass function (for which we usethe \citet{She99} parameterisation), and ${\rm d}r/{\rm
+d}z = c/H(z)$ with $H(z)=H_0 E(z) = H_0\sqrt{\Om(1+z)^3+\Ol}$ for a flat universe."
+ 1n order to carry out this integral. the power-law fits obtained in Section + can be used to represent pLGM.z) which give the mean relation between the halo mass and the DLA cross-section.," In order to carry out this integral, the power-law fits obtained in Section \ref{section:cross} can be used to represent $\sdla(M,z)$ which give the mean relation between the halo mass and the DLA cross-section."
+ Note that the dependence on the Llubble constant disappears on the right-hancl-sicle of equation (8) because dr£dz scales as h1. while naycd depends on h. and epa scales as h7 in the simulation.," Note that the dependence on the Hubble constant disappears on the right-hand-side of equation (8) because ${\rm d}r/{\rm d}z$ scales as $\hinv$, while $n_{\rm dm}{\rm d}M$ depends on $h^3$, and $\sdla$ scales as $h^{-2}$ in the simulation."
+ 1n Figure 5.. we show the cumulative abundance of DLAs per unit redshift at >=3 as a function of total halo mass.," In Figure \ref{cum_z3.eps}, we show the cumulative abundance of DLAs per unit redshift at $z=3$ as a function of total halo mass."
+ The horizontal shaded region in the left panel indicates the observed DLA abundance of Pérouxetal.(2001)., The horizontal shaded region in the left panel indicates the observed DLA abundance of \citet{Per01}.
+. We note that the data-set analysed by Pérouxetal.(2001). inelucles that of Storrie-Lombardi&Wolfe (2000)... and a similar value for the DLA abundance was also reported by Storric-(2000).. It is encouraging that the DLA abundances found in our simulations agree well with the observed range.," We note that the data-set analysed by \citet{Per01} includes that of \citet{Sto00}, , and a similar value for the DLA abundance was also reported by \citet{Sto00}.. It is encouraging that the DLA abundances found in our simulations agree well with the observed range."
+The precisev onieasured frequencies of solar oscillations provile us with a unique tool to probe the solar interior with sutiicicnt accuracy.,The precisely measured frequencies of solar oscillations provide us with a unique tool to probe the solar interior with sufficient accuracy.
+ These frequencies are primarily deteruiued x the dynamical quantities like sound speed. (Lens voor the adiabatie index of the solar material aud a primary inversion of the observed frequencies vields the sound. speec and deusitv profiles inside the Sun (Cough 1985: Cough Isosovichey 1990: Cough Thompson 1991: Dzicnubowski et al. 1991:," These frequencies are primarily determined by the dynamical quantities like sound speed, density or the adiabatic index of the solar material and a primary inversion of the observed frequencies yields the sound speed and density profiles inside the Sun (Gough \cite{dog85}; Gough Kosovichev \cite{dog90}; Gough Thompson \cite{dog91}; Dziembowski et al. \cite{dz94};"
+ Autia Basu 199la: Basu et al. 1996:, Antia Basu \cite{ab94a}; Basu et al. \cite{b96};
+ Ciough et al. 19963)., Gough et al. \cite{dog96}) ).
+ Ou the other hand. in order to incr the temperature and chemical composition profiles acditional assumptions regarding the input plysics are recwired (Shibahashi 1993:: Autia Chitre 1995: Shibahashi Takata 1996: Ikosovichev 1996)).," On the other hand, in order to infer the temperature and chemical composition profiles additional assumptions regarding the input physics are required (Shibahashi \cite{shi93}; ; Antia Chitre \cite{ac95}; Shibahashi Takata \cite{st96}; Kosovichev \cite{kos96}) )."
+ Thus. the equons of thermal equilibrium euable us to determine the teuperature aud bydrogen abundance profiles in the solar interior provided the opacities. equation of state auk nuclear energy. generation rates are prescribed.," Thus, the equations of thermal equilibrium enable us to determine the temperature and hydrogen abundance profiles in the solar interior provided the opacities, equation of state and nuclear energy generation rates are prescribed."
+ Although he primary mversious can vield the sound speed to an accuracy of. the opacities aud nuclear reactioji ates are hardly known to comparable acctirmaev aud COlLISCGIieutlv. more systematic errors are introduced in tjese secondary mversous for teniperature and chemical composition.," Although the primary inversions can yield the sound speed to an accuracy of, the opacities and nuclear reaction rates are hardly known to comparable accuracy and consequently, more systematic errors are introduced in these secondary inversions for temperature and chemical composition."
+ There aro 2a 1uber of approaches adopted for secondary Versions., There are a number of approaches adopted for secondary inversions.
+ Ikosovichev (1996)) has eumploved the equations of theαπλα] equilibrium to express the lauges in primary variables (p.D4) iu terms of those oel secoidarv variabless (YZ) aud obtained equations connecting the freque1ο cifferences to variatiows in] vbunidaice profiles.," Kosovichev \cite{kos96}) ) has employed the equations of thermal equilibrium to express the changes in primary variables $\rho,\Gamma_1$ ) in terms of those in secondary variables $Y,Z$ ) and obtained equations connecting the frequency differences to variations in abundance profiles."
+ It should be noted that modifications oei Z profile mainly affect the opacities iu the solar oeterior while the equation of state and nuclear energy ecnoration rates are aflected to a much lesser extent., It should be noted that modifications in $Z$ profile mainly affect the opacities in the solar interior while the equation of state and nuclear energy generation rates are affected to a much lesser extent.
+ Such a procedure is essentially equivalent to fiudiug the Y profile aloie with the unecessary opacity modifications., Such a procedure is essentially equivalent to finding the $Y$ profile along with the necessary opacity modifications.
+ Shibahashi aud Takata (1996.. hereinafter ST96) adopt je standard opacities and nuclear reactiou rates o obtain re feniperaure anc chenical abuudauce profies usus 16 inverted sound speed xofile.," Shibahashi and Takata \cite{st96}, hereinafter ST96) adopt the standard opacities and nuclear reaction rates to obtain the temperature and chemical abundance profiles using the inverted sound speed profile."
+ Antia aucL Chitre (19€)5.. 1996)) set out to estimate ιο central eniperature of the Sun.," Antia and Chitre \cite{ac95}, \cite{ac96}) ) set out to estimate the central temperature of the Sun."
+ Thev adoated the oeweyed souid speed aux density profiles to oltain the enperature CZ) and οτι abundance (Y) profiles iu ie solar core. but the main difference was that opacity aud melear reaction rates were not directly «3uploved or tus purpose.," They adopted the inverted sound speed and density profiles to obtain the temperature $T$ ) and helium abundance $Y$ ) profiles in the solar core, but the main difference was that opacity and nuclear reaction rates were not directly employed for this purpose."
+ Tustead. the T aud Y profiles were obtained by minimizing the varialon I opacities from the Stauard values.," Instead, the $T$ and $Y$ profiles were obtained by minimizing the variation in opacities from the standard values."
+ The main reasou for allowing variations in theoretically deteriuuedquatity like opacity rather, The main reason for allowing variations in theoretically determinedquantity like opacity rather
+"(2001).. Mane Tor an old population has variations with metallicity ~2x that of Mj,gc. an effect confirmed in (he empirical investigation of Pielrzviskietal.(2010).. who looked at the RC of 15 nearby galaxies observed with LST.",", $M_{V,RC}$ for an old population has variations with metallicity $\sim$ $\times$ that of $M_{I,RC}$, an effect confirmed in the empirical investigation of \citet{2010AJ....140.1038P}, who looked at the RC of 15 nearby galaxies observed with HST."
+ Moreover. any residual differenGal reddening will be ~2x as significant in V.," Moreover, any residual differential reddening will be $\sim$ $\times$ as significant in $V$."
+ These (wo effects render the bulge RC non-horizontal in V. [further complicating the fitting routine.," These two effects render the bulge RC non-horizontal in $V$, further complicating the fitting routine."
+ There is a difference between (his parameter and that predominantly used in the literature., There is a difference between this parameter and that predominantly used in the literature.
+" We computed AJ#CS"" whereas most results present AWPU72.", We computed ${\Delta}I^{RGBB}_{RC}$ whereas most results present ${\Delta}V^{RGBB}_{ZAHB}$.
+ However. these two values should be very nearly equal as the two largest biases are not large ancl go in opposing directions.," However, these two values should be very nearly equal as the two largest biases are not large and go in opposing directions."
+ Firstly. AWM? should be a little larger than ALCS because the RC will be a little bluer than the RG stars. the bias is expected [rom stellar theory. but is consistent with negligible in our data.," Firstly, ${\Delta}V^{RGBB}_{RC}$ should be a little larger than ${\Delta}I^{RGBB}_{RC}$ because the RC will be a little bluer than the RG stars, the bias is expected from stellar theory but is consistent with negligible in our data."
+" Conversely. AV}(7% will be a little smaller than AVE?"" as the ZAIB is the dimmest phase of horizontal branch evolution. however in our analvsis of elobular cluster data we find that this effect can be no more than 70.05 mag."," Conversely, ${\Delta}V^{RGBB}_{ZAHB}$ will be a little smaller than ${\Delta}V^{RGBB}_{RC}$ as the ZAHB is the dimmest phase of horizontal branch evolution, however in our analysis of globular cluster data we find that this effect can be no more than $\sim$ 0.05 mag."
+" We thus adopt (he approximation A[8655R¢_=""ROBLESAZ.", We thus adopt the approximation ${\Delta}I^{RGBB}_{RC} = {\Delta}V^{RGBB}_{ZAHB}$.
+" It is reassuring that the Galactic bulge has the faintest RGBB relative to its horizontal branch as it is the mostmetal-rich RGBB detected thus far,", It is reassuring that the Galactic bulge has the faintest RGBB relative to its horizontal branch as it is the mostmetal-rich RGBB detected thus far.
+ Galactic elobular clusters do not typically reach metallicities as high as |M/ILI] 0.0. and those that do have substantial differential reddening (Ortolanietal.2001) or multiple stellar populations 2009).. effects that render the RGBB harcler to detect.," Galactic globular clusters do not typically reach metallicities as high as [M/H] $\approx 0.0$, and those that do have substantial differential reddening \citep{2001A&A...376..878O} or multiple stellar populations \citep{2009Natur.462..483F}, effects that render the RGBB harder to detect."
+ This expansion of the parameter space al the metal rich end follows recent. complementary detections of the RGBB in vounger and moremetal-poor svstenms. those of the nearby cwarl galaxies.," This expansion of the parameter space at the metal rich end follows recent, complementary detections of the RGBB in younger and moremetal-poor systems, those of the nearby dwarf galaxies."
+ There have been detections of the RGDD toward the Sculptor dwarl spheroidal galaxy (Majewskietal.1999).. Ursa Minor (Bellazzinietal.2002).. Sagittarius (Monacoetal.2002).. and Sextans (Leeetal.2003).. and M32 (Monachesietal. 2011).. Mon," There have been detections of the RGBB toward the Sculptor dwarf spheroidal galaxy \citep{1999ApJ...520L..33M}, Ursa Minor \citep{2002AJ....124.3222B}, Sagittarius \citep{2002ApJ...578L..47M}, and Sextans \citep{2003AJ....126.2840L}, and M32 \citep{2011ApJ...727...55M}. ."
+ellietal.(2010) recently reported on the detection of the RGBB toward Cetus. 101010. LGS 3 and Tucana.," \citet{2010ApJ...718..707M} recently reported on the detection of the RGBB toward Cetus, IC1613, LGS 3 and Tucana."
+ The observed relation between AVore and [M/IT] is shown in Figure 3.. with the Zinn&West(1984). metallicity scale asstuned for the Galactic globular clusters and w C'en.," The observed relation between ${\Delta}V^{RGBB}_{ZAHB}$ and [M/H] is shown in Figure \ref{Fig:BumpEmpiricalHistory}, with the \citet{1984ApJS...55...45Z} metallicity scale assumed for the Galactic globular clusters and $\omega$ $Cen$."
+ We estimate [M/II] via the conversion function suggested by Salarisetal. (1993): and an [a /Fe]=+04 for the Galactic globular clusters ance Cen. and [a /Fe]=+0.25 for the Galactic bulge ancl M32.," We estimate [M/H] via the conversion function suggested by \citet{1993ApJ...414..580S}: : and an $\alpha$ $=+0.4$ for the Galactic globular clusters and$\omega$ $Cen$ , and $\alpha$ $=+0.25$ for the Galactic bulge and M32."
+ , 
+emission as well.,emission as well.
+ This suggests two possible results for the method's application to real data., This suggests two possible results for the method's application to real data.
+ First. if there are few or no soft sources just slightly harder than the canonical SSSs. the procedure will select only SSSs.," First, if there are few or no soft sources just slightly harder than the canonical SSSs, the procedure will select only SSSs."
+ Second. if there is a supply of somewhat harder sources. we will select them as well.," Second, if there is a supply of somewhat harder sources, we will select them as well."
+ We have now had opportunities to apply the selection criteria to the 4+ galaxies studied in the companion paper Kong 2003). to M104 eet al.," We have now had opportunities to apply the selection criteria to the $4$ galaxies studied in the companion paper Kong 2003), to M104 et al."
+ 2003a). M31 eet al.," 2003a), M31 et al."
+ 2003b). and roughly one dozen additional. galaxies eet al.," 2003b), and roughly one dozen additional galaxies et al."
+ 2003c)., 2003c).
+ We have found that most galaxies have significant populations of both SSSs and. sources with somewhat harder spectra (e.g.. ΚΑΤ«250 eV).," We have found that most galaxies have significant populations of both SSSs and sources with somewhat harder spectra (e.g., $k\, T < 250$ eV)."
+ We refer to the latter as quasisoft sources (QSSs)., We refer to the latter as quasisoft sources (QSSs).
+ Our selection procedure distinguishes between SSSs and QSSs according to which step in the algorithm identifies the source as being very soft., Our selection procedure distinguishes between SSSs and QSSs according to which step in the algorithm identifies the source as being very soft.
+ In the galaxies we have studied. spectral fits for the brightest SSS and QSS candidates have verified that the algorithmic classification works.," In the galaxies we have studied, spectral fits for the brightest SSS and QSS candidates have verified that the algorithmic classification works."
+ To simplify the terminology. we will sometimes use the term “very soft source” (VSS) to refer to both SSSs and QSSs.," To simplify the terminology, we will sometimes use the term “very soft source"" (VSS) to refer to both SSSs and QSSs."
+ The physical significance of this new class is not yet understood. but there are likely to be several physical models corresponding to QSSs.," The physical significance of this new class is not yet understood, but there are likely to be several physical models corresponding to QSSs."
+ First. hot SSSs located behind large gas columns will have photons in the medium energy band. M (1.122 keV). but may have few photons in the soft band. S (0.1—1.1 keV).," First, hot SSSs located behind large gas columns will have photons in the medium energy band, $M$ $1.1-2$ keV), but may have few photons in the soft band, $S$ $0.1-1.1$ keV)."
+ For such sources. the hardness ratios typically used to identify SSSs will have values not normally associated with SSSs. even though their intrinsic characteristics clearly place them in the SSS category.," For such sources, the hardness ratios typically used to identify SSSs will have values not normally associated with SSSs, even though their intrinsic characteristics clearly place them in the SSS category."
+ Second if the detector has poorer than anticipated sensitivity to soft photons. soft sources can appear to be harder than they actually are.," Second if the detector has poorer than anticipated sensitivity to soft photons, soft sources can appear to be harder than they actually are."
+ Thus. some QSSs are likely to have the same physical characteristics as some other sources identified as SSSs.," Thus, some QSSs are likely to have the same physical characteristics as some other sources identified as SSSs."
+ Finally. some QSSs are likely to be genuinely harder than SSSs. so hard that white dwarf models can be ruled out.," Finally, some QSSs are likely to be genuinely harder than SSSs, so hard that white dwarf models can be ruled out."
+ As we will discuss in 82. intermediate mass black hole models may be appropriate for such systems. but neutron star or stellar mass black hole models should also be considered.," As we will discuss in 2, intermediate mass black hole models may be appropriate for such systems, but neutron star or stellar mass black hole models should also be considered."
+ Below we list some of the questions we hope to answer with studies that compare VSS populations in different galaxies. (, Below we list some of the questions we hope to answer with studies that compare VSS populations in different galaxies. (
+1) What are typical galactic populations of SSSs and QSSs?,1) What are typical galactic populations of SSSs and QSSs?
+ Irrespective of their fundamental natures. the answer to this question will allow us to estimate the influence of soft X-ray sources as Ionizers of the ISM. (," Irrespective of their fundamental natures, the answer to this question will allow us to estimate the influence of soft X-ray sources as ionizers of the ISM. ("
+2) Are any spiral galaxy parameters related to the relative sizes of SSS and QSS populations?,2) Are any spiral galaxy parameters related to the relative sizes of SSS and QSS populations?
+ Answering this question can provide insight into the age of the populations that spawn very soft sources. and hence might help to illuminate their nature.," Answering this question can provide insight into the age of the populations that spawn very soft sources, and hence might help to illuminate their nature."
+ Rappaport (1994) suggested that for spiral galaxies. the size of the SSS population might scale with blue luminosity. but this has not been tested. (," Rappaport (1994) suggested that for spiral galaxies, the size of the SSS population might scale with blue luminosity, but this has not been tested. ("
+3) Do elliptical galaxies house large SSS/QSS populations?,3) Do elliptical galaxies house large SSS/QSS populations?
+ Although it has been suggested that the diffuse soft emission in ellipticals may be due to SSSs (see. e.g.. Fabbiano. Kim. Trinchiert 1994). we still know very little about SSSs in ellipticals.," Although it has been suggested that the diffuse soft emission in ellipticals may be due to SSSs (see, e.g., Fabbiano, Kim, Trinchieri 1994), we still know very little about SSSs in ellipticals."
+ If accreting WDs form the largest segment of SSS populations. and if a significant fraction of the donor stars have masses small enough to be typical of the stars found in elliptical galaxies. then we may expect SSSs to be important parts of the X-ray source population in ellipticals. (," If accreting WDs form the largest segment of SSS populations, and if a significant fraction of the donor stars have masses small enough to be typical of the stars found in elliptical galaxies, then we may expect SSSs to be important parts of the X-ray source population in ellipticals. ("
+4) Within spiral galaxies. what are the relative populations of SSSs/QSSs in the galaxy bulges and disks? (,"4) Within spiral galaxies, what are the relative populations of SSSs/QSSs in the galaxy bulges and disks? ("
+5) Do galaxies with massive central black holes have more SSSs or QSSs located within | kpe of the nucleus than comparable galaxies without massive central black holes?,5) Do galaxies with massive central black holes have more SSSs or QSSs located within 1 kpc of the nucleus than comparable galaxies without massive central black holes?
+ It has been suggested that some SSSs within the central kpc of galaxies which harbor massive black holes may actually be the stripped cores of stars that have been tidally disrupted eet al., It has been suggested that some SSSs within the central kpc of galaxies which harbor massive black holes may actually be the stripped cores of stars that have been tidally disrupted et al.
+ 2001)., 2001).
+ Verification of this hypothesis by studying individual SSSs will be difficult. so statistical studies of SSS populations in a large number of galaxies may provide the best tests. (," Verification of this hypothesis by studying individual SSSs will be difficult, so statistical studies of SSS populations in a large number of galaxies may provide the best tests. ("
+6) For all galaxies. are the positions of QSSs and SSSs correlated to the positions of other objects. such as HII regions. planetary nebulae. supernova remnants. or globular clusters?,"6) For all galaxies, are the positions of QSSs and SSSs correlated to the positions of other objects, such as HII regions, planetary nebulae, supernova remnants, or globular clusters?"
+ The distances to most external galaxies are too large to allow for convincing optical identifications., The distances to most external galaxies are too large to allow for convincing optical identifications.
+ It is nevertheless useful to identify the types of populations which tend to be associated with VSSs., It is nevertheless useful to identify the types of populations which tend to be associated with VSSs.
+ This can provide clues to their fundamental natures. (, This can provide clues to their fundamental natures. (
+7) Are SSSs significant contributors to the rates of Type la SNe?,7) Are SSSs significant contributors to the rates of Type Ia SNe?
+ The answer to this question can. be achieved. by combining information about typical total galactic populations with studies of the viability of the accreting WD models., The answer to this question can be achieved by combining information about typical total galactic populations with studies of the viability of the accreting WD models.
+ Previous studies of SSSs in external galaxies have used a variety of selection criteria., Previous studies of SSSs in external galaxies have used a variety of selection criteria.
+" In NGC 4697, e.g.. Sarazin. Irwin. Bregman (2001) identified 3 SSSs by requiring that HRI=(M—S)/(M+S)=-1 and HR2=(H—S)/(H+S)=—1. where S. M. and H represent the numbers of counts in the bands 0.3—| keV. 1-2 keV. and 2—IO keV. In their studies of the colors of X-ray sources. Prestwich et (2002). used the same criteria, which are satisfied by only a handful of the sources they analyzed. drawn from both MIOI and M83."," In NGC 4697, e.g., Sarazin, Irwin, Bregman (2001) identified $3$ SSSs by requiring that ${\tilde {HR1}}=(\tilde M-\tilde S)/(\tilde M+\tilde S) = -1$ and ${\tilde {HR2}}=(\tilde H-\tilde S)/(\tilde H+\tilde S) = -1,$ where $\tilde S,$ $\tilde M,$ and $\tilde H$ represent the numbers of counts in the bands $0.3-1$ keV, $1-2$ keV, and $2-10$ keV, In their studies of the colors of X-ray sources, Prestwich et (2002), used the same criteria, which are satisfied by only a handful of the sources they analyzed, drawn from both M101 and M83."
+ Less restrictive criteria were used by Swartz et (2002) to identify SSSs in M81., Less restrictive criteria were used by Swartz et (2002) to identify SSSs in M81.
+ The criteria HRI<—0.5. HR2——0.5 selected 12 M8] X-ray sources. 2 of which were eliminated because they are identified with foreground stars. while one is identified with a supernova remnant (SNR).," The criteria ${\tilde {HR1}} < -0.5,$ ${\tilde {HR2}} < -0.5$ selected $12$ M81 X-ray sources, $2$ of which were eliminated because they are identified with foreground stars, while one is identified with a supernova remnant (SNR)."
+ Pence et (2002a) identified 10 SSSs in MIOI. but did not specify the selection criteria.," Pence et (2002a) identified $10$ SSSs in M101, but did not specify the selection criteria."
+ The possible physical interpretation of the sources seemed to play a role. as one of the galaxy’s softest sources was not counted among the SSSs. perhaps because tt appears to be too luminous to be a nuclear-burning WD (Pence et 2002b).," The possible physical interpretation of the sources seemed to play a role, as one of the galaxy's softest sources was not counted among the SSSs, perhaps because it appears to be too luminous to be a nuclear-burning WD (Pence et 2002b)."
+" Kong et (2002a) took another approach. requiring (HR240gsx—]and (HRI <0.) CHRI0,4 <—0.8])."," Kong et (2002a) took another approach, requiring $({\tilde {HR2}} + \sigma_{{\tilde {HR2}}} \leq -1$ $[{\tilde {HR1}} < 0,$ ] ${\tilde {HR1}} + \sigma_{{\tilde {HR1}}} \leq -0.8]$ )."
+ Fourteen sources in the central 17«17! of M31 satisfied these conditions. of which 2 were apparently identified with SNRs (Kong et al.," Fourteen sources in the central $17' \times 17'$ of M31 satisfied these conditions, of which $2$ were apparently identified with SNRs (Kong et al."
+ 2002b) and 3 with possible foreground stars., 2002b) and $3$ with possible foreground stars.
+ These latter criteria were developed in parallel with a study, These latter criteria were developed in parallel with a study
+the models.,the models.
+ The DFs of the best models for FS373 and FS76 are presented in Figs., The DFs of the best models for FS373 and FS76 are presented in Figs.
+ 12. and 13.. respectively.," \ref{mod373} and \ref{mod76}, respectively."
+ We plot the DF in the equatorial plane in turning-point space., We plot the DF in the equatorial plane in turning-point space.
+" Each orbit in this plane is labeled uniquely by its pericenter distance Aa; and apocenter distance Ray, if. A4; 18 given the same sign as L..", Each orbit in this plane is labeled uniquely by its pericenter distance $R_{\rm peri}$ and apocenter distance $R_{\rm apo}$ if $R_{\rm peri}$ is given the same sign as $L_z$.
+ Circular orbits lie on two straight lines with Rayo=ER., Circular orbits lie on two straight lines with $R_{\rm apo} = \pm R_{\rm peri}$.
+ Radial orbits lie on the vertical line with Ανν=0., Radial orbits lie on the vertical line with $R_{\rm peri}=0$.
+ In both galaxy models. an excess phase-space density of stars on near-circular orbits. forming the KDC. is clearly visible.," In both galaxy models, an excess phase-space density of stars on near-circular orbits, forming the KDC, is clearly visible."
+ Moreover. the KDC is obviously disjunct from the central nucleus or density cusp.," Moreover, the KDC is obviously disjunct from the central nucleus or density cusp."
+ Since the KDCs form a distinct subcomponent within their host galaxies. the stars that make up a KDC can be singled out of the DF and be studied separately (especially in the case of FS373. it was very clear which basis functions in the expansion of the DF formed the KDC).," Since the KDCs form a distinct subcomponent within their host galaxies, the stars that make up a KDC can be singled out of the DF and be studied separately (especially in the case of FS373, it was very clear which basis functions in the expansion of the DF formed the KDC)."
+ In order to roughly estimate the stellar mass of the KDC. we assumed a stellar mass-to-light ratio of Μην=2—AMLp. which agrees with the observed colors and line-strengths.," In order to roughly estimate the stellar mass of the KDC, we assumed a stellar mass-to-light ratio of $M/L_B = 2-4 M_\odot/L_{B,\odot}$, which agrees with the observed colors and line-strengths."
+ Thus. we find Μίκης=173x107M. for both galaxies or a few percent at most of the total mass.," Thus, we find $M_{\rm
+KDC} \approx 1-5 \times 10^7 M_\odot$ for both galaxies or a few percent at most of the total mass."
+ The adopted M/Ly is both typical fora 10 Gyr old. metal-poor (21 «[Fe/H]« —-0.5) stellar population (which would agree with dEs being primordial stellar systems) and for a 5 Gyr old. more metal-rich (-0.5. «|Fe/H]« 0.0) stellar population (which would agree with dEs being harassed late-type spirals that experienced a starburst) (Worthey (1994))).," The adopted $M/L_B$ is both typical for a 10 Gyr old, metal-poor $-1<$ $<-0.5$ ) stellar population (which would agree with dEs being primordial stellar systems) and for a 5 Gyr old, more metal-rich $-0.5<$ $<0.0$ ) stellar population (which would agree with dEs being harassed late-type spirals that experienced a starburst) \cite{wo94}) )."
+ The key question is whether KDCs in dwarf elliptical galaxies are produced the same way as in massive ellipticals., The key question is whether KDCs in dwarf elliptical galaxies are produced the same way as in massive ellipticals.
+ We explore two possible avenues to KDC formation in. dEs., We explore two possible avenues to KDC formation in dEs.
+ The first is the merger hypothesis. as in. giant. ellipticals: the second is the harassment scenario. which posits that. gravitational interactions play an important role in the evolution of dEs.," The first is the merger hypothesis, as in giant ellipticals; the second is the harassment scenario, which posits that gravitational interactions play an important role in the evolution of dEs."
+ The analytical arguments given below are strictly speaking only valid for anddistant encounters., The analytical arguments given below are strictly speaking only valid for and encounters.
+" An encounter between to galaxies. with nasses M, and M». qualifies as if. at closest approach. the change in the potential energy of the pair is much smaller than the initial orbital kinetic energy."," An encounter between to galaxies, with masses $M_1$ and $M_2$, qualifies as if, at closest approach, the change in the potential energy of the pair is much smaller than the initial orbital kinetic energy."
+ In afast encounter. the relative velocity of the galaxies is much larger than the internal stellar velocities.," In a encounter, the relative velocity of the galaxies is much larger than the internal stellar velocities."
+ This translates into the following constraints on the impact parameter 5 and the internal velocity dispersion cn: with V4.4 the relative velocity of the interacting galaxies., This translates into the following constraints on the impact parameter $b$ and the internal velocity dispersion $\sigma_{\rm int}$: with $V_{\rm rel}$ the relative velocity of the interacting galaxies.
+ For Ms9κΙΟΥM. a typical dE mass. and Mj=Mpc<<M. we find 6>250—500 pe for Vig=c300—400 km/s. Also. cp>Oi.," For $M_2 \approx 5 \times 10^9 M_\odot$, a typical dE mass, and $M_1 =
+M_{\rm KDC} << M$, we find $b > 250 - 500$ pc for $V_{\rm rel} =
+\sigma_{\rm gal} = 300 - 400$ km/s. Also, $\sigma_{\rm gal} >
+\sigma_{\rm int}$."
+ Hence. any non-penetrating encounter between a dE and a much smaller dwarf galaxy classifies as a fast and distant encounter (even if we take the dwarf galaxy to be originally 10 times more massive than Mypc. the minimum impact parameter would change by only 10%)).," Hence, any non-penetrating encounter between a dE and a much smaller dwarf galaxy classifies as a fast and distant encounter (even if we take the dwarf galaxy to be originally 10 times more massive than $M_{\rm KDC}$, the minimum impact parameter would change by only )."
+" In the case of a giant elliptical with M»=5κ10!'M. and M,=Ma<<M». the condition for a fast flyby becomes b>25—50 kpe. again rather unsensitive to Maji."," In the case of a giant elliptical with $M_2
+\approx 5 \times 10^{11} M_\odot$ and $M_1 = M_{\rm dE} << M_2$, the condition for a fast flyby becomes $b > 25 - 50$ kpc, again rather unsensitive to $M_{\rm dE}$."
+ In a group or cluster environment. galaxies keep respectable distances of a few tens of kpe (Mooreefaf.(1996))).," In a group or cluster environment, galaxies keep respectable distances of a few tens of kpc \cite{mkldo}) )."
+ With this in mind. we can discuss possible mechanism of producing KDCs in dEs.," With this in mind, we can discuss possible mechanism of producing KDCs in dEs."
+ While the merger origin of KDCs in bright ellipticals is well accepted. anumber of facts argue against the merger hypothesis in the case of dEs.," While the merger origin of KDCs in bright ellipticals is well accepted, a number of facts argue against the merger hypothesis in the case of dEs."
+" The change of the forward velocity of a galaxy with mass M, induced by a fast. distant hyperbolic encounter with a galaxy with mass M» with a relative velocity V4 is given by (Sparke&Gallagher(2000).. Binney&Tremaine (1987)))."," The change of the forward velocity of a galaxy with mass $M_1$ induced by a fast, distant hyperbolic encounter with a galaxy with mass $M_2$ with a relative velocity $V_{\rm rel}$ is given by \cite{sg},, \cite{bt}) )."
+. The closer and the slower the encounter. the more orbital energy is converted into internal (stellar) kinetic energy.," The closer and the slower the encounter, the more orbital energy is converted into internal (stellar) kinetic energy."
+ For an encounter between a typical M;=5xIOM. dE anda Mj=5x10’M. dwarf galaxy with a relative velocity V4=c300 km/s. AV) is very small (e.g. AV)~35 km/s for a collision with 6=1 kpe).," For an encounter between a typical $M_2 =5 \times 10^9
+M_\odot$ dE and a $M_1 =5 \times 10^7 M_\odot$ dwarf galaxy with a relative velocity $V_{\rm rel} = \sigma_{\rm gal} = 300$ km/s, $\Delta
+V_{||}$ is very small (e.g. $\Delta V_{||} \sim 35$ km/s for a collision with $b=1$ kpc)."
+" In the case of an encounter between a M»=5x10!'M, elliptical and a M;=5x10°M.. dE. on the other hand. the velocity change is substantial: AV)~Vier. even for impact parameters of a few tens of kiloparsecs."," In the case of an encounter between a $M_2=5 \times 10^{11} M_\odot$ elliptical and a $M_1=5 \times 10^9
+M_\odot$ dE, on the other hand, the velocity change is substantial: $\Delta V_{||} \sim V_{\rm rel}$, even for impact parameters of a few tens of kiloparsecs."
+ This suggests that a dE. in a group or cluster environment. has virtually no chance of slowing down and capturing another (smaller) dwarf galaxy. contrary to a more massive elliptical galaxy.," This suggests that a dE, in a group or cluster environment, has virtually no chance of slowing down and capturing another (smaller) dwarf galaxy, contrary to a more massive elliptical galaxy."
+ Hence. once the galaxy group or cluster is 1n place. the chance of forming a KDC in a dE by à merger is exceedingly small.," Hence, once the galaxy group or cluster is in place, the chance of forming a KDC in a dE by a merger is exceedingly small."
+ Also. it is unclear how the merger scenario can explain the complex velocity profile of FS373. particularly the velocity changing sign around a radius of 12”=2.4 kpe.," Also, it is unclear how the merger scenario can explain the complex velocity profile of FS373, particularly the velocity changing sign around a radius of $12'' = 2.4$ kpc."
+ Alternatively. the merger could have taken place the group or cluster virialized. in an environment where relative velocities were smaller than the present values.," Alternatively, the merger could have taken place the group or cluster virialized, in an environment where relative velocities were smaller than the present values."
+ The low galaxy density in such an environment argues against this idea., The low galaxy density in such an environment argues against this idea.
+ Also. it remains to be seen. e.g. using high-resolution N- simulations. whether à KDC formed this way can survive the dE’s falling into à group or cluster and the subsequent gravitational interactions with giant group or cluster members.," Also, it remains to be seen, e.g. using high-resolution $N$ -body simulations, whether a KDC formed this way can survive the dE's falling into a group or cluster and the subsequent gravitational interactions with giant group or cluster members."
+ A plausible alternative is the spin-up of a dE's halo by fly-by encounters with other galaxies., A plausible alternative is the spin-up of a dE's halo by fly-by encounters with other galaxies.
+" The impulse approximation and the tensor virial theorem yield the following expression for the maximum amount of angular momentum that can be transfered to a galaxy with mass M, during an encounter with a galaxy with mass M»:with qi the axis ratio and{τι a component of the inertial tensor (S", The impulse approximation and the tensor virial theorem yield the following expression for the maximum amount of angular momentum that can be transfered to a galaxy with mass $M_1$ during an encounter with a galaxy with mass $M_2$ :with $q_1$ the axis ratio and$I_{11}$ a component of the inertial tensor \cite{ssk}) ).
+omSunder&Kochhar (1990))). Using AJ.~ to roughly estimate Άγιοι the maximum possible," Using $\Delta J \sim M_1 R_{\rm e,1} \Delta
+v_{\rm rot}$ to roughly estimate $\Delta v_{\rm rot}$ , the maximum possible"
+data on the Sn rates and star formation rates to infer the relative role played by type Ia and II Sn (SnIa and SnII hereafter).,data on the Sn rates and star formation rates to infer the relative role played by type Ia and II Sn (SnIa and SnII hereafter).
+ A different approach was pursued by other authors. which considered. gas-dynamical mechanisms that at relatively low redshift are responsible for redistributing previously produced metals.," A different approach was pursued by other authors, which considered gas–dynamical mechanisms that at relatively low redshift are responsible for redistributing previously produced metals."
+ For instance. ? suggested that clumps of low-entropy highly enriched gas may sink in the central cluster regions. thereby leading to an increase of the observed emission—weighted metallicity.," For instance, \cite{2008arXiv0802.0975C} suggested that clumps of low–entropy highly enriched gas may sink in the central cluster regions, thereby leading to an increase of the observed emission–weighted metallicity."
+ For instance. ram—pressure stripping of the interstellar medium (SM) of merging galaxies has been suggested as a mechanism o pollute at relatively low redshift a metal-poor ICM with highly enriched gas (e.g.2.andreferencestherein). while causing a morphological transformation of cluster galaxies (e.g.22).," For instance, ram–pressure stripping of the interstellar medium (ISM) of merging galaxies has been suggested as a mechanism to pollute at relatively low redshift a metal–poor ICM with highly enriched gas \citep[e.g.,][and references
+therein]{2006A&A...452..795D}, while causing a morphological transformation of cluster galaxies \citep[e.g.,][]{2007MNRAS.tmpL..43C,2007MNRAS.380.1399R}."
+ Although possible evidences of ram-pressure stripping of cluster galaxies have been detected (e.z..?) the question remains as to whether this mechanism dominates the evolution of the ICM enrichment.," Although possible evidences of ram–pressure stripping of cluster galaxies have been detected \citep[e.g.,][]{2007ApJ...659L.115C} the question remains as to whether this mechanism dominates the evolution of the ICM enrichment."
+ Indeed. since ram pressure is expected to be more efficient in high-temperature clusters. one expects an increasing trend of metallicity with ICM temperature (e.g.2y..," Indeed, since ram pressure is expected to be more efficient in high–temperature clusters, one expects an increasing trend of metallicity with ICM temperature \citep[e.g.,][]{1997ApJ...488...35R}."
+ If any. observations suggest that hotter systems have a relatively lower metallicity (e.g..2).. thus suggesting that ram-pressure stripping is not the dominant process in enriching the ICM.," If any, observations suggest that hotter systems have a relatively lower metallicity \citep[e.g.,][]{2005ApJ...620..680B}, thus suggesting that ram–pressure stripping is not the dominant process in enriching the ICM."
+ It is clear that understanding the history of the ICM enrichment in cosmological context. during the cluster hierarchical build up. requires describing in detail the gasdynamics related to the merging processes. while including a self-consistent treatment of star formation and chemical evolution.," It is clear that understanding the history of the ICM enrichment in cosmological context, during the cluster hierarchical build up, requires describing in detail the gasdynamics related to the merging processes, while including a self–consistent treatment of star formation and chemical evolution."
+ In this context. cosmological hydrodynamical simulations offer a unique means to capture in full detail the complexity of these processes (e.g. 2222. see ?.. for a recent review).," In this context, cosmological hydrodynamical simulations offer a unique means to capture in full detail the complexity of these processes (e.g., \citealt{2003MNRAS.339.1117V,2004MNRAS.349L..19T,2006MNRAS.371..548R,2007MNRAS.382.1050T}, see \citealt{2008arXiv0801.1062B}, for a recent review)."
+ In their mos advanced versions. chemo-dynamical simulation codes treat the production of different metal species. released by different stellar populations by resorting to detailed stellar yields. also accounting for the dependent stellar lifetimes.," In their most advanced versions, chemo–dynamical simulation codes treat the production of different metal species, released by different stellar populations by resorting to detailed stellar yields, also accounting for the mass--dependent stellar lifetimes."
+ In this paper we will present results on the ICM metal abundance from cosmological simulations of galaxy clusters. using the chemo-dynamical version of the ccode (??1. which has been recently presented by ? (TO7 hereafter).," In this paper we will present results on the ICM metal abundance from cosmological simulations of galaxy clusters, using the chemo–dynamical version of the code \citep{SP01.1,2005MNRAS.364.1105S}, which has been recently presented by \cite{2007MNRAS.382.1050T} (T07 hereafter)."
+ We will compare the simulations with observational— results on the Iron abundance profiles.Zi... of nearby clusters. on the evolution of the ICM metallicity and on the SnIa rates.," We will compare the simulations with observational results on the Iron abundance profiles, of nearby clusters, on the evolution of the ICM metallicity and on the SnIa rates."
+ This comparison will be performed with the aim of shading light on the relative role played by star formation. feedback processes and gus dynamics in determining the cosmic history of metal enrichment.," This comparison will be performed with the aim of shading light on the relative role played by star formation, feedback processes and gas dynamics in determining the cosmic history of metal enrichment."
+ The plan of the paper is as follows., The plan of the paper is as follows.
+ In Section 2 we review our implementation of chemical evolution in the ccode and present the main characteristics of the cluster simulations., In Section 2 we review our implementation of chemical evolution in the code and present the main characteristics of the cluster simulations.
+ Section 3 will be devoted to the comparison between simulation results and observations., Section 3 will be devoted to the comparison between simulation results and observations.
+ After comparing the profiles of the Iron abundance. we will concentrate on the evolution of the ICM metallicity.," After comparing the profiles of the Iron abundance, we will concentrate on the evolution of the ICM metallicity."
+ We will then compare observations and simulation predictions on the rate of SnIu., We will then compare observations and simulation predictions on the rate of SnIa.
+ We will draw our conclusions in Section 4., We will draw our conclusions in Section 4.
+ All values of Tron abundance that we will quote in the following are scaled to the solar abundance value by ?.., All values of Iron abundance that we will quote in the following are scaled to the solar abundance value by \cite{1998SSRv...85..161G}.
+" In this letter we present a set of simulations of four massive isolated clusters. which have been identified in a Dark-Matter only simulation having a box size 47941Mpe (2)... performed for a flat ACDM cosmological model with €,=0.3. Pius=0.7. c4=0.9 and O,=0.04."," In this letter we present a set of simulations of four massive isolated clusters, which have been identified in a Dark–Matter only simulation having a box size $479 \hm$ \citep{2001MNRAS.328..669Y}, performed for a flat $\Lambda$ CDM cosmological model with $\Omega_m =
+0.3$, $h_{100} =0.7$ , $\sigma_8 = 0.9$ and $\Omega_b = 0.04$."
+ The four extracted Lagrangian regions. centred on these clusters with virial in the range Alas—1.022.3 107ΑΙ. have been resimulated using the Zoomed Initial Condition (ZIC) technique by ?.. which allows one to increase force and mass resolution in the regions of interest.," The four extracted Lagrangian regions, centred on these clusters with virial in the range $M_{\rm
+ vir}=$ $\,\times 10^{15} \msun$, have been resimulated using the Zoomed Initial Condition (ZIC) technique by \cite{TO97.2}, which allows one to increase force and mass resolution in the regions of interest."
+" The high-resolution DM particles have mass mpi;=1.1810°PhAL... and the barionie particles have been added with a mass my,=L7105b.FAL. inorderto reproduce the assumec cosmic barionic fraction."," The high–resolution DM particles have mass $m_{DM}=1.13\times 10^9 \msun$ , and the barionic particles have been added with a mass $m_{gas}=1.7\times 10^8 \msun$ in order to reproduce the assumed cosmic barionic fraction."
+ The basic characteristics of the simulated clusters are summarized in Table I.., The basic characteristics of the simulated clusters are summarized in Table \ref{tab:simul}. .
+ The simulations are performed using the hydrodynamica Tree-SPH code (9). with the implementation of chemical enrichment by TO7., The simulations are performed using the hydrodynamical Tree-SPH code \citep{2005MNRAS.364.1105S} with the implementation of chemical enrichment by T07.
+ The Plummer-equivalent softening length for gravitational force is set to ο=5htkpe in physica units from ο=2 to >=(0. while at higher redshifts is ο=15htkpe in comoving units.," The Plummer–equivalent softening length for gravitational force is set to $\epsilon = 5 \hk$ in physical units from $z=2$ to $z=0$, while at higher redshifts is $\epsilon = 15 \hk$ in comoving units."
+ The simulations include heating from a uniform time-dependent UV background (2). anc metallicity-dependent radiative cooling based on the tables by for an optically thin plasma., The simulations include heating from a uniform time-dependent UV background \citep{1996ApJ...461...20H} and metallicity–dependent radiative cooling based on the tables by \cite{1993ApJS...88..253S} for an optically thin plasma.
+ The process of star formation (SF hereafter) is described by the sub-resolution multiphase model by ?.. for which the density threshold for the onset of SF is set to ny=O1lem ," The process of star formation (SF hereafter) is described by the sub–resolution multiphase model by \cite{2003MNRAS.339..289S}, for which the density threshold for the onset of SF is set to $n_H=0.1\,$ $^{-3}$."
+ While the relevant features of the chemical evolution model are described here below. we address the reader to TO7 for a more detailed description.," While the relevant features of the chemical evolution model are described here below, we address the reader to T07 for a more detailed description."
+ Metals are produced by SnII. SnIa and intermediate and low-mass stars (ILMS hereafter). with only SnIa and SnII providing energy feedback.," Metals are produced by SnII, SnIa and intermediate and low–mass stars (ILMS hereafter), with only SnIa and SnII providing energy feedback."
+ We assume SnII to arise from stars having mass above SA/.., We assume SnII to arise from stars having mass above $8M_\odot$.
+" As for the SnIa. we assume their progenitors to be binary systems. whose total mass lies in the range (3-16,U.."," As for the SnIa, we assume their progenitors to be binary systems, whose total mass lies in the range $M_\odot$."
+ Metals and energy are released by stars of different mass by properly accounting for mass—dependent lifetimes., Metals and energy are released by stars of different mass by properly accounting for mass–dependent lifetimes.
+ In this work we assume the lifetime function proposed by ?.., In this work we assume the lifetime function proposed by \cite{1993ApJ...416...26Pb}.
+ We adopt the metallicity-dependent stellar yields by 3 for SnIIL. the yields by ? for the ILMS and by ? for SnIa.," We adopt the metallicity–dependent stellar yields by \cite{1995ApJS..101..181W} for SnII, the yields by \cite{1997A&AS..123..305V} for the ILMS and by \cite{2003NuPhA.718..139T} for SnIa."
+ The version of the code used for the simulations presented here allowed us to follow H. He. C. N. ο. Mz. Si and Fe.," The version of the code used for the simulations presented here allowed us to follow H, He, C, N, O, Mg, Si and Fe."
+ Once produced by a star particle. metals are then spread to the surrounding gas particles by using the B-spline kerne with weights comouted over 64 neighbours and taken to be proportional to the volume of each particle.," Once produced by a star particle, metals are then spread to the surrounding gas particles by using the B-spline kernel with weights computed over 64 neighbours and taken to be proportional to the volume of each particle."
+ TO7 verified with detailed tests that the final results on the pattern of chemical enrichment are ratyer insensitive to tle Weighting scheme (kernel shape andnumber of neighbours) used to spread metals., T07 verified with detailed tests that the final results on the pattern of chemical enrichment are rather insensitive to the weighting scheme (kernel shape andnumber of neighbours) used to spread metals.
+ Our simulations include the kinetic feedback model implemented by ?» ., Our simulations include the kinetic feedback model implemented by \cite{2003MNRAS.339..289S}. .
+ According to tus scheme. SnII explosions trigger galactic winds. whose mass upload rate is assumed to be proportional to the star formation rae. Aly=WAL.," According to this scheme, SnII explosions trigger galactic winds, whose mass upload rate is assumed to be proportional to the star formation rate, $\dot{M}_W =\eta
+\dot{M}_{\star}$."
+ Therefore. fixing the parameter 7) and the wind velocity ey; amounts to fix the total energy carried by the winds.," Therefore, fixing the parameter $\eta$ and the wind velocity $v_W$ amounts to fix the total energy carried by the winds."
+ Our choice of 7—3 and my=500kms1 corresponds to assume.for the initial mass function (IMF) by ?.. with SnIIreleasing 10'1 ergs each. nearly unity efficiency in powering galactic outflows.," Our choice of $\eta=3$ and $v_W = 500\vel$ corresponds to assume,for the initial mass function (IMF) by \cite{1955ApJ...121..161S}, , with SnIIreleasing $10^{51}$ ergs each, nearly unity efficiency in powering galactic outflows."
+ In our comparison with observational data. we will first explore the effect of changing the IMF.," In our comparison with observational data, we will first explore the effect of changing the IMF."
+ We use the IMF by 9? and, We use the IMF by \cite{1955ApJ...121..161S} and
+in Fig. 105. ,"in Fig. \ref{fig:m_hi_optsize}) ),"
+although with a larger scatter., although with a larger scatter.
+ A linear fit with a slope and intercept of 1.744-0.2? and 6.93-0.18. respectively is shown as a solid line.," A linear fit with a slope and intercept of $\pm$ 0.22 and $\pm$ 0.18, respectively is shown as a solid line."
+ The larger scatter in the relation between Mu and the optical diameter. also seen in sample of brighter dwarfs (e.g. Swaters (1999))). is probably indicative of a looser coupling between the gas and star formation in dwarfs. compared to that in spiral galaxies.," The larger scatter in the relation between ${\rm{_{HI}}}$ and the optical diameter, also seen in sample of brighter dwarfs (e.g. \cite{swater99}) ), is probably indicative of a looser coupling between the gas and star formation in dwarfs, compared to that in spiral galaxies."
+ Figure 11. shows the HI mass to light ratio for the FIGGS sample plotted asa function of the HI extent. Dui/Dy...," Figure \ref{fig:mtol_size} shows the HI mass to light ratio for the FIGGS sample plotted asa function of the HI extent, ${\rm{_{HI}/D_{Ho}}}$."
+ A trend of an increase in the My)/Lis with an increase in the HI extent of the galaxies is clearly seen., A trend of an increase in the ${\rm{_{HI}/L_B}}$ with an increase in the HI extent of the galaxies is clearly seen.
+ The best fit to the FIGGS sample shown as a solid line gives van Zee et al.(1995) from a HI mapping of a sample of low luminosity galaxies also found an evidence of an extended HI extent for high Mui/Ly galaxies., The best fit to the FIGGS sample shown as a solid line gives van Zee et al.(1995) from a HI mapping of a sample of low luminosity galaxies also found an evidence of an extended HI extent for high ${\rm{_{HI}/L_B}}$ galaxies.
+ Figure 12. shows Myij/Li for the FIGGS sample as a function of My., Figure \ref{fig:mtol_lb} shows ${\rm{_{HI}/L_B}}$ for the FIGGS sample as a function of ${\rm{_B}}$ .
+ The same quantity for several other spiral and, The same quantity for several other spiral and
+infinitesimally thin. relativistic self-gravitating disces with internal pressure.,"infinitesimally thin, relativistic self-gravitating discs with internal pressure."
+ The pressure is given by a polytropic equation of state., The pressure is given by a polytropic equation of state.
+ The. polvtropic exponent +=3 was used. since in this case there exists an cxact solution in the ανομία limit.," The polytropic exponent $\gamma=3$ was used, since in this case there exists an exact solution in the Newtonian limit."
+ This special choice corresponds to three-dimensional bodies of constant densities as well (see LEunter 1972)., This special choice corresponds to three-dimensional bodies of constant densities as well (see Hunter 1972).
+ Thus. we may make a comparison of the results ounce here for Hat disces with those of rotating. homogeneous relativistic stars.," Thus, we may make a comparison of the results found here for flat discs with those of rotating, homogeneous relativistic stars."
+ As Butterworth Ipser (1976) showed. sequences of homogeneous. rotating relativistic stars usually. erminate at the mass-shee limit.," As Butterworth Ipser (1976) showed, sequences of homogeneous, rotating relativistic stars usually terminate at the mass-shed limit."
+ Close to the Newtonian case however. they were not able to follow the sequences o this limit.," Close to the Newtonian case however, they were not able to follow the sequences to this limit."
+ We suggest that. as weaker relativistic disces Xfurcate into a ring. weaker relativistic stars of constant clensity bifurcate into a toroid structure as well. ancl do not end at the mass-shed limit.," We suggest that, as weaker relativistic discs bifurcate into a ring, weaker relativistic stars of constant density bifurcate into a toroid structure as well, and do not end at the mass-shed limit."
+ In the Newtonian limit. the structure of these constant density toroids were calculated by Eriguchi Sugimoto (1981).," In the Newtonian limit, the structure of these constant density toroids were calculated by Eriguchi Sugimoto (1981)."
+ The pressure-less disces possess ergo-regions. where the drageing of inertial frames would force observers to rotate.," The pressure-less discs possess ergo-regions, where the dragging of inertial frames would force observers to rotate."
+ These first appear at a single point within the disc at =[4l and reach the edge of the disc at z=1.89 (Meine IxIeinwacchter 1993)., These first appear at a single point within the disc at $z\subscr{e}=1.41$ and reach the edge of the disc at $z=1.89$ (Meinel Kleinwäcchter 1993).
+ Phev also occur in rotating stars (Butterworth Lpser 1976)., They also occur in rotating stars (Butterworth Ipser 1976).
+ For disces with internal oressure however. we found. no indication of the existence of ergo-regions.," For discs with internal pressure however, we found no indication of the existence of ergo-regions."
+ It is to be expected. that there exists a continuous transition from the zero-pressure A=Q0 case o disces with non-vanishing pressure (A. small)., It is to be expected that there exists a continuous transition from the zero-pressure $K=0$ case to discs with non-vanishing pressure $K$ small).
+ However. his connection could not be demonstrated numerically. as disc sequences end either at the mass-shecl limit or bifurcate into rings before reaching the pressure-less limit.," However, this connection could not be demonstrated numerically, as disc sequences end either at the mass-shed limit or bifurcate into rings before reaching the pressure-less limit."
+ One might speculate. perhaps. that such a connection can be achieved by a disc which consists of two parts: a pressure supported central region (0«p py). surrounded. by a dust. disc (pympx pa).," One might speculate, perhaps, that such a connection can be achieved by a disc which consists of two parts: a pressure supported central region $0 < \rho < \rho\subscr{p}$ ), surrounded by a dust disc $\rho\subscr{p} < \rho < \rho\subscr{d}$ )."
+" Phe dust cise would refer to pj,=0. and the disc sequences to py=pa."," The dust disc would refer to $\rho\subscr{p}=0$, and the disc sequences to $\rho\subscr{p}=\rho\subscr{d}$."
+ In some rotating relativistic stars. there are sequences of supra-massive stars (Cook. Shapiro “Teukolsky 1992). which are so massive that they exceed the rest mass of a non-rotating star. and can only exist for non-zero rotation.," In some rotating relativistic stars, there are sequences of supra-massive stars (Cook, Shapiro Teukolsky 1992), which are so massive that they exceed the rest mass of a non-rotating star, and can only exist for non-zero rotation."
+ In the case of Dat. relativistic discs with non-zero internal pressure. we did not find any supra-massive disc sequences.," In the case of flat relativistic discs with non-zero internal pressure, we did not find any supra-massive disc sequences."
+ Possibly. these negative results are a result of the particular equation of state where in the strong relativistic limit the mass ds always concentrated in the centre of the disc. as shown in Fig.," Possibly, these negative results are a result of the particular equation of state where in the strong relativistic limit the mass is always concentrated in the centre of the disc, as shown in Fig."
+ 4 for the non-rotating case., 4 for the non-rotating case.
+ Other relations between pressure and surface density. in particular a smaller ὃν could. also possibly leac το supra-niassivo sequences or ergo-reglons.," Other relations between pressure and surface density, in particular a smaller $\gamma$, could also possibly lead to supra-massive sequences or ergo-regions."
+ The numerical method. developed here is. sullicientLy general to be applied to the study of cifferentially rotating discs. and can casily be extended. to. threc-cdimensional rotating stars as well.," The numerical method developed here is sufficiently general to be applied to the study of differentially rotating discs, and can easily be extended to three-dimensional rotating stars as well."
+ This may be the subject of a future further investigation., This may be the subject of a future further investigation.
+ A stability analysis of the computed equilibrium. configurations lies bevonc the scope of the present papor., A stability analysis of the computed equilibrium configurations lies beyond the scope of the present paper.
+ Useful discussions with Drs it. Aleinel and. T. Wolf are eratefullys acknowledged., Useful discussions with Drs R. Meinel and T. Wolf are gratefully acknowledged.
+including RR Lyrae stars. TRGB. and Cepheids.,"including RR Lyrae stars, TRGB, and Cepheids."
+ Insofar as all these methods start from a common. fundamental Local Group distance seale that has been carefully assembled by a combination of many standard candles. the ~£0.2- mag variances noted above are reflective of the small discrepancies that can emerge from the different methods in ways that are often still hard to pinpoint and that are not even consistent from one galaxy to another.," Insofar as all these methods start from a common, fundamental Local Group distance scale that has been carefully assembled by a combination of many standard candles, the $\sim \pm0.2-$ mag variances noted above are reflective of the small discrepancies that can emerge from the different methods in ways that are often still hard to pinpoint and that are not even consistent from one galaxy to another."
+ For a more extensive discussion on another system (NGC 5128) where several stellar candles can be accurately compared. see Harris et al. (2010)).," For a more extensive discussion on another system (NGC 5128) where several stellar candles can be accurately compared, see Harris et al. \cite{har10}) )."
+ Another direct. though somewhat less precise. comparison method of interest is the linear size distribution of globular clusters (GCs). which has been developed by Jordánn et ((2005)) into a standard-ruler technique.," Another direct, though somewhat less precise, comparison method of interest is the linear size distribution of globular clusters (GCs), which has been developed by Jordánn et \cite{jor05}) ) into a standard-ruler technique."
+ The key quantity is the peak of the GC half-light radius distribution. normalized to host galaxy size and calibrated via the Milky Way GCs.," The key quantity is the peak of the GC half-light radius distribution, normalized to host galaxy size and calibrated via the Milky Way GCs."
+ Using the Jordann et al., Using the Jordánn et al.
+ data for M87 and their calibration. we obtain d=(16.442.3) Mpe or Ga—M);=31.07+0.30.," data for M87 and their calibration, we obtain $d = (16.4 \pm
+2.3)$ Mpc or $(m-M)_0 = 31.07 \pm 0.30$."
+" Combining the four methods listed above (TRGB. PNLF. GC sizes. SBF/Cepheids). we obtain a weighted average distance modulus (n—M),31.08+0.06 for M87. or d(16.4+0.5) Mpe."," Combining the four methods listed above (TRGB, PNLF, GC sizes, SBF/Cepheids), we obtain a weighted average distance modulus $\langle m-M \rangle_0 = 31.08
+\pm 0.06$ for M87, or $d = (16.4 \pm 0.5)$ Mpc."
+ A more precise TRGB distance especially could be obtained very straightforwardly with halo-star photometry in a less crowded region. and would in our view be the most effective way to calibrate the distance to this important galaxy.," A more precise TRGB distance especially could be obtained very straightforwardly with halo-star photometry in a less crowded region, and would in our view be the most effective way to calibrate the distance to this important galaxy."
+ To test the internal errors and completeness of the photometry we ran two separate artificial-star procedures., To test the internal errors and completeness of the photometry we ran two separate artificial-star procedures.
+ Inthe first series. stars were added to a representative 1500κ 500-px region of the image in the lower left corner (similar to the region shown in Figure 1).," Inthe first series, stars were added to a representative $1500 \times 500-$ px region of the image in the lower left corner (similar to the region shown in Figure 1)."
+ These added stars were distributed evenly in color and magnitude over the ranges 0.5«(V—7)5.0 and 27.0«I« 29.6. as shown in Figure 7..," These added stars were distributed evenly in color and magnitude over the ranges $0.5 < (V-I) < 5.0$ and $27.0 < I < 29.6$ , as shown in Figure \ref{cmd_fake}."
+ These intervals deliberately covered a larger range in both color and magnitude than in our observed CMD (Figure 3))., These intervals deliberately covered a larger range in both color and magnitude than in our observed CMD (Figure \ref{cmd4}) ).
+ The same measurement sequence as done on the original frames was then carried out. with a two-pass sequence off," The same measurement sequence as done on the original frames was then carried out, with a two-pass sequence of."
+ind/phot/allstar.. Figure 7. shows the results for inserted stars that were actually recovered in the photometry. including both their input magnitudes and colors (center panel) and their actually measured values (right panel).," Figure \ref{cmd_fake} shows the results for inserted stars that were actually recovered in the photometry, including both their input magnitudes and colors (center panel) and their actually measured values (right panel)."
+ Of the total of 10000 added stars. just 1917 of these were successfully found and measured in both colors. and a high fraction of these lie in the upper left (bright. blue) part of the CMD.," Of the total of 10000 added stars, just 1917 of these were successfully found and measured in both colors, and a high fraction of these lie in the upper left (bright, blue) part of the CMD."
+ The completeness of detection fniCrecovered)/n(input) as a function of magnitude is shown in Figure 8:: the completeness levels are reached at FOOOW=29.00 and F814W=28.15.," The completeness of detection $f =
+n(recovered)/n(input)$ as a function of magnitude is shown in Figure \ref{completeness}; the completeness levels are reached at $F606W = 29.00$ and $F814W = 28.15$."
+" To be classified as ""recovered"" a star must be detected in both images.", To be classified as “recovered” a star must be detected in both images.
+ Note that at very faint levels (below the point) the formal values of f tend to decrease rather slowly. a result of the very high degree of crowding.," Note that at very faint levels (below the point) the formal values of $f$ tend to decrease rather slowly, a result of the very high degree of crowding."
+ Figure 7 (particularly the difference between the second and third panels) clearly indicates that internal random uncertainties are large at all levels of the CMD., Figure \ref{cmd_fake} (particularly the difference between the second and third panels) clearly indicates that internal random uncertainties are large at all levels of the CMD.
+ Figure 9 displays the differences between the measured and input magnitudes more completely., Figure \ref{random} displays the differences between the measured and input magnitudes more completely.
+ For Jz28.0. one magnitude below the RGB tip. completeness of detection becomes low and the systematic errors increase. such that most stars are measured too bright.," For $I \gtrsim 28.0$, one magnitude below the RGB tip, completeness of detection becomes low and the systematic errors increase, such that most stars are measured too bright."
+ We do not consider this faint region further., We do not consider this faint region further.
+ The measured are on average slightly too blue independent of magnitude. by (ACV—7))=—0.22 mag.," The measured are on average slightly too blue independent of magnitude, by $\langle \Delta (V-I)
+\rangle = -0.22$ mag."
+ Over the top magnitude of the RGB (727— 28) the internal random scatter of the magnitudes and colors (rms) is 7=£0.36 mag in F814W. £0.53 mag in F606W. and £0.40 mag in (F606W— F814W).," Over the top magnitude of the RGB $I = 27 - 28$ ) the internal random scatter of the magnitudes and colors (rms) is $\sigma =
+\pm0.36$ mag in $F814W$, $\pm0.53$ mag in $F606W$, and $\pm 0.40$ mag in $(F606W-F814W)$ ."
+ The artificially even distribution of input stars in Figure 7 overpopulates the brightest part of the luminosity function compared with the real CMD.," The artificially even distribution of input stars in Figure \ref{cmd_fake}
+ overpopulates the brightest part of the luminosity function compared with the real CMD."
+ A second run of artificial-star, A second run of artificial-star
+large (7)..,large \citep{1998T}.
+" The parazuueters of this model are Neptune's ""Uóutial (defined below) senuüauajor axis. ecceutricitv and inclination. the planets migration rate. and the timescales for Neptune eccentricity and inclination damping."," The parameters of this model are Neptune's “initial"" (defined below) semi-major axis, eccentricity and inclination, the planet's migration rate, and the timescales for Neptune's eccentricity and inclination damping."
+ In this model. we include the effects of only oue plauet (Neptune). an approach we justify briefly in Section 3.3. and more thoroughly in ?..," In this model, we include the effects of only one planet (Neptune), an approach we justify briefly in Section \ref{subsec:justnep} and more thoroughly in \citet{2012D}."
+ We define Noptuue's orbital evolution using the following parameters: ere we consider what range of parameters we should explore for Neptunes migration direction. distance. aud timescale: Direct computational modeling of the effect of planctesimals on Neptune’s orbit would be colputationally expensive. so iusteac we apply fictitious forces (Appendix) to evolve Neptune’s soimianajor axis ay. eccentricity ον. and inclination xw. with anv specified functional form.," We define Neptune's orbital evolution using the following parameters: Here we consider what range of parameters we should explore for Neptune's migration direction, distance, and timescale: Direct computational modeling of the effect of planetesimals on Neptune's orbit would be computationally expensive, so instead we apply fictitious forces (Appendix) to evolve Neptune's semi-major axis $a_N$, eccentricity $e_N$, and inclination $i_N$, with any specified functional form."
+ Following ?— and ?.. we use the functional forms: where ey is the initial senuiuajor axis of Neptune. d; = 30 AU is the final semi-major axis. and 7. 7;. aud Τι ave the eccentricity damping timescale. inclination damping timescale. aud migration timescale respectively.," Following \citet{1993M} and \citet{2008L}, we use the functional forms: where $a_0$ is the initial semi-major axis of Neptune, $a_f$ = 30 AU is the final semi-major axis, and $\tau_e$, $\tau_i$, and $\tau_a$ are the eccentricity damping timescale, inclination damping timescale, and migration timescale respectively."
+ Oi results do not depeud ou the specific form of Equ (1))., Our results do not depend on the specific form of Eqn \ref{eqn:forms}) ).
+" As we will demonstrate in Section L.. sometimes the instantaneous rate of change of the variables 5.ipds "" most relevant. while m other cases the total(4, evolution,matters most."," As we will demonstrate in Section \ref{sec:results}, sometimes the instantaneous rate of change of the variables $\frac{\dot{a}}{a}, \frac{\dot{e}}{e}, \frac{\dot{i}}{i}$ ) is most relevant, while in other cases the total evolutionmatters most."
+ We have verified these statements wit1 integratious (not shown) using an alternative mieratio1 fori 46«x;= constaut., We have verified these statements with integrations (not shown) using an alternative migration form $\frac{\dot{a}}{a} \varpropto \frac{\dot{e}}{e} \varpropto \frac{\dot{i}}{i} \equiv {\rm constant}$ .
+ We model an initially unexcited disk of planctesimals that becomes todaws cold classical population., We model an initially unexcited disk of planetesimals that becomes today's cold classical population.
+ Iu Section 3.1... we present the observational coustraiuts ou the excitation of this population.," In Section \ref{sec:obs}, we present the observational constraints on the excitation of this population."
+ Iu Section 3.2.. we present an analytical model for the evolution of this planetesimal disk uuder the influence of Neptune. which we use to predict and interpret the results of umnucrical simulations.," In Section \ref{subsec:sec}, we present an analytical model for the evolution of this planetesimal disk under the influence of Neptune, which we use to predict and interpret the results of numerical simulations."
+ Iu Section 3.3.. we justify directly modeling ouly Neptune instead of all four giant plaucts.," In Section \ref{subsec:justnep}, we justify directly modeling only Neptune instead of all four giant planets."
+" The cold classicals are a class of dyviiuuicallv ""cold? objects on low-ceceutricity. low-inclination orbits. with positions starting at 12.5 AU. the reeion interior to which 15is unstable. aud falling off quickly bevoud 15 AU (?).."," The cold classicals are a class of dynamically “cold"" objects on low-eccentricity, low-inclination orbits, with positions starting at 42.5 AU, the region interior to which is unstable, and falling off quickly beyond 45 AU \citep{2009K}."
+ We assune that todav's cold classical IKRBOs are remnant plauctesimals that formed sins and we use theseternis interchangeably.," We assume that today's cold classical KBOs are remnant planetesimals that formed , and we use theseterms interchangeably."
+ Stroug constraints cau be placed, Strong constraints can be placed
+The observable part. of debris disks are small (: lumn) clusty or icy grains. collisionally produce [rom larger. undetectable parent bodies.,"The observable part of debris disks are small $\leq 1\,$ mm) dusty or icy grains, collisionally produced from larger, undetectable parent bodies."
+ In addition to the gravitational »ull of the star. these grains are also alfected by several orces such as stellar radiation pressure. Povnting-Robertson (DIU) drag and the possible gravitational inlluence of large »odies in the neighborhood.," In addition to the gravitational pull of the star, these grains are also affected by several forces such as stellar radiation pressure, Poynting-Robertson (PR) drag and the possible gravitational influence of large bodies in the neighborhood."
+ As has been shown in numerous numerical studies. the combined. effect of hese different orces can lead to complex spatial structures in resolved disks. (c.g.ναι2008).," As has been shown in numerous numerical studies, the combined effect of these different forces can lead to complex spatial structures in resolved disks \citep[e.g.][]{wyatt08}."
+. A less investigated: additional orce that could. have an influence on grain dynamics is he crag due to particles [rom the surrounding interstellar medium (1981)., A less investigated additional force that could have an influence on grain dynamics is the drag due to particles from the surrounding interstellar medium (ISM).
+ Phe elfect of ISM has first been addressed w Artvmowicz&Clampin(1997).. who studied the level of disk erosion. due to sandblasting by ISAT dust. @rains.," The effect of ISM has first been addressed by \citet{arty97}, who studied the level of disk erosion due to sandblasting by ISM dust grains."
+ They coneluded that. at least around: massive stars. this ellect was negligible because small LSAT grains felt a strong repulsive radiation force.," They concluded that, at least around massive stars, this effect was negligible because small ISM grains felt a strong repulsive radiation force."
+ More recentlv.Scherer. (2000).. Debesetal.(2009)... Manessctal. (2009)... Belvaey&Ralikov(2010) and Pastor(2011)| considered. instead the οσοι of ISAL on disk grains.," More \citet{scherer}, , \citet{debes09}, \citet{manes09}, \citet{bera} and \citet{pasto} considered instead the effect of ISM on disk grains."
+ Vhis Πακ of neutral atoms acts indeed similarly to the solar wind or radiation pressure from a physical point of view but. being monocdirectional. can significantly perturb the trajectories of the grains. ancl potentially induce asvnimetric structures in the disk.," This flux of neutral atoms acts indeed similarly to the solar wind or radiation pressure from a physical point of view but, being monodirectional, can significantly perturb the trajectories of the grains, and potentially induce asymmetric structures in the disk."
+ In particular Manessctal.(2009). ancl Debesetal.(2009) sugeest that the ISM. Bux. can explain the unusual morphology of some debris disks like HD61005 anc 1132997., In particular \citet{manes09} and \citet{debes09} suggest that the ISM flux can explain the unusual morphology of some debris disks like HD61005 and HD32997.
+ In their model Debesetal.(2009). consider dus particles close to the blow-out size for the star and compute the trajectories of perturbed grains over a timescale of 5000 ves., In their model \citet{debes09} consider dust particles close to the blow-out size for the star and compute the trajectories of perturbed grains over a timescale of 5000 yrs.
+ The majority of their grains are strongly perturbed ane enc up quickly on hyperbolic orbits., The majority of their grains are strongly perturbed and end up quickly on hyperbolic orbits.
+ A similar scenario is outlined by Manessetal.(2009) where they concentrate on small grains (0.1 pam) whose lifetime before ejection is of the order of afew 10°? vears., A similar scenario is outlined by \citet{manes09} where they concentrate on small grains (0.1 $\mu$ m) whose lifetime before ejection is of the order of afew $10^3$ years.
+ The morphology changes they, The morphology changes they
+"to LO ke utιο MECS 1.55 to 10 keV: the IIPCGSPC 7 to 65 keV axd the PDS 13 to 200 keV. To ensure that he fittius statisic D (A7) was unbiased across the eutire cnerey range. data were rebinuned to a width of of the ""ull width at ial :anaxinun cherev resolution function of ecli lustrewt and also under the condition that each iu contaiue ΠΕ of 20 counts.","to 10 keV; the MECS 1.85 to 10 keV; the HPGSPC 7 to 65 keV and the PDS 13 to 200 keV. To ensure that the fitting statistic $\chi^2$ ) was unbiased across the entire energy range, data were rebinnned to a width of of the full width at half maximum energy resolution function of each instrument and also under the condition that each bin contained a minimum of 20 counts."
+ It is known from iuter-lustreit spectral calibrations that there can be simall position ¢lepeudeut normalization differences between he lustreits., It is known from inter-instrument spectral calibrations that there can be small position dependent normalization differences between the instruments.
+ Therefore. these factors were included as free multiplicative parameters during multiple NFI spectral fittine.," Therefore, these factors were included as free multiplicative parameters during multiple NFI spectral fitting."
+ Fie., Fig.
+ ls rows low aud high cucerey light curves measired by the MECS (1.310.0 keV) aud the PDS (€3200) NV) after background subtraction., 1 shows low and high energy light curves measured by the MECS (1.3–10.0 keV) and the PDS (13–200 keV) after background subtraction.
+ The time resolulon dis [00 s. The sotrco intensity is clearly variable with a doubliug of the fiux occuring on the time scale of adjacent biis., The time resolution is 400 s. The source intensity is clearly variable with a doubling of the flux occurring on the time scale of adjacent bins.
+ The fastest observed fluctuation in the MECS is on a time scale of —10 s. consistent with the ASCA resul of Kubo e al. (," The fastest observed fluctuation in the MECS is on a time scale of $\sim$ 10 s, consistent with the ASCA result of Kubo et al. ("
+1998).,1998).
+ In the PDS. the fastes observable variatioLis of the order of  10 s. This time πεale implies an upper iuit on the size of the cussion region of a few. « 10H. cx.," In the PDS, the fastest observable variation is of the order of $\sim$ 40 s. This time scale implies an upper limit on the size of the emission region of a few $\times$ $^{11}$ cm."
+ Next a search for periodic variations was carried out., Next a search for periodic variations was carried out.
+ On short time scales (fντα ! 0.5 Iz). a powcr density curve reveals à l/f tvpo distribution. |mt no clear periodicities (see Fie.," On short time scales $f \sim$ $^{-4}$ –0.5 Hz), a power density curve reveals a ${1/f}$ –type distribution, but no clear periodicities (see Fig."
+ 2)., 2).
+ At ΠΙΟ longer periocls. a period search reveals weak chhancements at 35 aud 115 nius.," At much longer periods, a period search reveals weak enhancements at 35 and 145 mins."
+ We estimate hat for narrow QPO/periodicities iu the 0.01 lz &>» 0.1 Tz range. we could detect a zuuplitude moclation at the 3o level.," We estimate that for narrow QPO/periodicities in the 0.01 Hz to 0.1 Hz range, we could detect a amplitude modulation at the $\sigma$ level."
+" Iu order to studv the loug-terii X-ray vaability of ~ Cas, we have also analyzed R-NTE All Sky Monitor (ASN) data which has coutinuously observed the source from 1996 February. 20 to 1998 Deceuber. 31."," In order to study the long-term X-ray variability of $\gamma$ –Cas, we have also analyzed R-XTE All Sky Monitor (ASM) data which has continuously observed the source from 1996 February, 20 to 1998 December, 31."
+ We have searched for periods in the range 30?500 days using the Lomb-Scarele periodoeram. using oth the individual dwell aud i-dav averages data.," We have searched for periods in the range 30–500 days using the Lomb-Scargle periodogram, using both the individual dwell and 1-day averages data."
+ No cals wdhoa lig[um sjeuilficance (e... >99%)) were ford in the individual dwel data. although one peak at a periocL of —20 davs was of mareinal significance (a the e»t) level).," No peaks with a high significance (i.e., $>$ ) were found in the individual dwell data, although one peak at a period of $\sim$ 200 days was of marginal significance (at the $\sim$ level)."
+" As a check. «un dts reality, we replace the R-XTE ASA measurements with data drawn from <Vv Gassa cüstribution ceitered on zero With ao of 1 (i.e. 10 signal)."," As a check on its reality, we replaced the R-XTE ASM measurements with data drawn from a Gaussian distribution centered on zero with a $\sigma$ of 1 (i.e. no signal)."
+ We note that he peak still existed (albei with a lucu lower significance). implying that it is probably caused by a windowing cifect and therefore does not reflect a true period in the N-vay flux of y Cas.," We note that the peak still existed (albeit with a much lower significance), implying that it is probably caused by a windowing effect and therefore does not reflect a true period in the X-ray flux of $\gamma$ –Cas."
+ Analysis of the l-dav averaged data vielded. sinilu results. i.c.. no sieuificaut periodicities.," Analysis of the 1-day averaged data yielded similar results, i.e., no significant periodicities."
+ We estunate that for periods aroun 100 davs. we would have detected a periodic modulation at the confidence level.," We estimate that for periods around 100 days, we would have detected a periodic modulation at the confidence level."
+ Siuple models (ie. power-laws. bremsstrahlung. ete.)," Simple models (i.e., power-laws, bremsstrahlung, etc.)"
+ eave poor fits to the data., gave poor fits to the data.
+ For example. a1. absorbed power-law gives a X of GAL for 203 degrees of freedom (dof).," For example, an absorbed power-law gives a $\chi ^2$ of 684 for 203 degrees of freedom (dof)."
+ The addition of au ion line at 0.03 keV iuproves the fit sienificautly (\?/dof=505/2WO). but is still unacceptable at energies above I) keV. Based ou previous ROSAT aud ASC'A measuremeits CIaberl 1995: Ίνπνο ct al.," The addition of an iron line at $\pm$ 0.03 keV improves the fit significantly $\chi ^2$ /dof=505/200), but is still unacceptable at energies above 10 keV. Based on previous ROSAT and ASCA measurements (Haberl 1995; Kubo et al."
+ 1998). we nest investigated opically thin thermal plasina models with both the eniperature aud elemental abundance as free parameters.," 1998), we next investigated optically thin thermal plasma models with both the temperature and elemental abundance as free parameters."
+ The results are listed in Table 1., The results are listed in Table 1.
+ A best-fit MEIAL model. based. iu," A best-fit MEKAL model, based in"
+the disc illuiunatiou is not taken iuto account.,the disc illumination is not taken into account.
+ Whereas ina nonirradiated disc small amplitude outbursts appear when the heating frout cannot bring the whole disc iuto a hot state. here the disc never returus to quiescence iu its inner parts. so the cooling wave is reflected into a heating wave when it ects close to the stable hot inner part of the disc.," Whereas in a non–irradiated disc small amplitude outbursts appear when the heating front cannot bring the whole disc into a hot state, here the disc never returns to quiescence in its inner parts, so the cooling wave is reflected into a heating wave when it gets close to the stable hot inner part of the disc."
+ The amplitude of these reflares grows as a consequence of the culanced mass transfer durius asin. until the disc mass has grown up to a point where a self-sustained long outburst is possible.," The amplitude of these reflares grows as a consequence of the enhanced mass transfer during maximum, until the disc mass has grown up to a point where a self-sustained long outburst is possible."
+ The refiaves properties also depend on the ratio μονρα2 the smaller this ratio. the more important he veflares (see Menou et al.," The reflares properties also depend on the ratio $\alpha_{\rm hot}/\alpha_{\rm
+cold}$: the smaller this ratio, the more important the reflares (see Menou et al."
+ 1999. for a discussion of lis effect in the context of XN-rav tranusicuts)., \cite{mhln99} for a discussion of this effect in the context of X-ray transients).
+ This is simply due to the fact hat. the lower this ratio. the larger MNaas after the passage of a cooling frout. Mayas boiug he maximum surface deusitv on the cold stable brauch: in the limiting case μοι=65444 there are no outbursts (Sinals 198 1)). but a heating/coolne wave that propagates jack and forth.," This is simply due to the fact that, the lower this ratio, the larger $\Sigma / \Sigma_{\rm max}$ after the passage of a cooling front, $\Sigma_{\rm max}$ being the maximum surface density on the cold stable branch; in the limiting case $\alpha_{\rm hot} = \alpha_{\rm cold}$ there are no outbursts (Smak \cite{s84}) ), but a heating/cooling wave that propagates back and forth."
+ Fie., Fig.
+ 5 shows the effect of chaugiug oc o a παω: (by a factor 2) value.," \ref{fig:alpha} shows the effect of changing $\alpha_{\rm
+cold}$ to a smaller (by a factor 2) value."
+ Successive reflares no ouecr reach the outer edge of the disc. aud their amplitude herefore decreases from one imini-outbirst to the uext onc.," Successive reflares no longer reach the outer edge of the disc, and their amplitude therefore decreases from one mini-outburst to the next one."
+ This accounts for the presence of flat top outbursts which were absent in Fie., This accounts for the presence of flat top outbursts which were absent in Fig.
+" Lob. The lightcurves are similar or all mass transfer rates. showing the pattern of Fie 5 with longer recurrence times for smaller AÁ4,."," \ref{fig:mdot}b b. The lightcurves are similar for all mass transfer rates, showing the pattern of Fig \ref{fig:alpha}
+ with longer recurrence times for smaller $\dot{M}_{\rm tr}$."
+ The oulv exception is for 7«Lot! gs +t awhich is close to stability. and for which the main outbursts are of he outsicde-in type.," The only exception is for $7 \times 10^{16}$ g $^{-1}$, which is close to stability, and for which the main outbursts are of the outside-in type."
+ Since small discs favour large reflares. it is not surprisiug that when oue considers discs with average rou1059 cni. and one takes api = 0.02 and Ayor = 0.2. one obtains a combination of the lighteurves shown in the two previous sections: Fig 6 is a good example of this.," Since small discs favour large reflares, it is not surprising that when one considers discs with average $r_{\rm out} = 1.3 \times 10^{10}$ cm, and one takes $\alpha_{\rm cold}$ = 0.02 and $\alpha_{\rm hot}$ = 0.2, one obtains a combination of the lightcurves shown in the two previous sections; Fig \ref{fig:rout} is a good example of this."
+ It is worth notius that such a light curve is reniuisceut of that ofCue.. even though the timescales are not quite the same.," It is worth noting that such a light curve is reminiscent of that of, even though the timescales are not quite the same."
+ We do obtain the right pattern for the reflares. but we do not reproduce the very loug superoutburst of (100 days). that would require 5 to be very close to unitv. caning that the linear approximation in Eq.," We do obtain the right pattern for the reflares, but we do not reproduce the very long superoutburst of (100 days), that would require $\gamma$ to be very close to unity, meaning that the linear approximation in Eq."
+ l ds invalid., \ref{eq:ill_sec} is invalid.
+ ForSee.. one already. had to assmue a relatively large value of + (0.87) in order to reproduce the observed 25 days duration: since the outburst duration varies as l/log(2) (Wameury et al. 1997)).," For, one already had to assume a relatively large value of $\gamma$ (0.87) in order to reproduce the observed 25 days duration; since the outburst duration varies as $1/\log(\gamma)$ (Hameury et al. \cite{hlh97}) ),"
+ we would need y=0.97 to obtain 100 davs., we would need $\gamma = 0.97$ to obtain 100 days.
+ Another differeuce with EG Cue is the amplitude of the 10ànjoutbursts: the observed ones have approximately the same amplitude. whereas we ect two identical nmuünioutbursts. the others being of decreasing amplitude.," Another difference with EG Cnc is the amplitude of the minioutbursts: the observed ones have approximately the same amplitude, whereas we get two identical minioutbursts, the others being of decreasing amplitude."
+ We have not been able to reproduce this behaviour with our parameterization: a possible solution is to introduce a time dependeut temperature of the white dwarf., We have not been able to reproduce this behaviour with our parameterization; a possible solution is to introduce a time dependent temperature of the white dwarf.
+ This is expected. because the superoutburst asted lone enough to heat up the surface of the white dwarf that will then cool.," This is expected, because the superoutburst lasted long enough to heat up the surface of the white dwarf that will then cool."
+" If the rebrightcnmes of ire dudeed due to illumination effects. his iuiplies that à5,44 cannot be πια] as we do not obtain reflares when ecg Is significantly less than 0.0! Osaki et al. (1997))"," If the rebrightenings of are indeed due to illumination effects, this implies that $\alpha_{\rm cold}$ cannot be small as we do not obtain reflares when $\alpha_{\rm cold}$ is significantly less than 0.01; Osaki et al. \cite{ost97}) )"
+ reached the same conclusion. but on different grounds: they assuued that Aeold Was increased to 0.1 caving the superoutburst. renamed hieh for 2 months. aud then decreased back to sunall values (0.001). and had therefore to set αμα to be an explicit fiction of time.," reached the same conclusion, but on different grounds; they assumed that $\alpha_{\rm cold}$ was increased to 0.1 during the superoutburst, remained high for 2 months, and then decreased back to small values (0.001), and had therefore to set $\alpha_{\rm cold}$ to be an explicit function of time."
+Aloute Carlo analyses suggest a marginal significauce of ~99% for this possible periodicity. though the low wuplitude aud siguificauce. aud the brief evcle-count of the 3.51 observations. demand additional data to verity.,"Monte Carlo analyses suggest a marginal significance of $\sim99\%$ for this possible periodicity, though the low amplitude and significance, and the brief cycle-count of the 3.5m observations, demand additional data to verify."
+ For SDSS JO092613621. both the l-hour timespan March sequence with 10 sec exposures (aud 60) sec time-resolution). and the longer 2-hour timespan April sequence with 20 sec exposures (and [0 κας time-resolution) reveal a strongly detected 28.3140.01 1uinute modulation. with striking eclipses (Figure 2).," For SDSS J0926+3624, both the 1-hour timespan March sequence with 40 sec exposures (and 60 sec time-resolution), and the longer 2-hour timespan April sequence with 20 sec exposures (and 40 sec time-resolution) reveal a strongly detected $\pm$ 0.01 minute modulation, with striking eclipses (Figure 3)."
+ From our Apri data. we estimate eclipse ceuters at IJD(TT)=2153173.725393(3)|E« 0.01966(1).," From our April data, we estimate eclipse centers at ${\rm HJD(TT)} = 2453473.725393(3) + E\times0.01966(1)$ ."
+ The eclipses are sharp and deep. with mid-eclipse depth of (at least) about Linagnitude.," The eclipses are sharp and deep, with mid-eclipse depth of (at least) about 1 magnitude."
+", A Caussian fit to the sharp portion of the eclipse duration vields a fullwidth at half iiiuinimua of about LO seconds and. for example. a ~1.3 münute duratiou full width at of the niuiuinua eclipse depth."," A Gaussian fit to the sharp portion of the eclipse duration yields a full-width at half minimum of about 40 seconds and, for example, a $\sim1.3$ minute duration full width at of the minimum eclipse depth."
+ As our ιοτοσοΊο is oulv of order LO sec. the eclipses could be even sharper/deeper.," As our time-resolution is only of order 40 sec, the eclipses could be even sharper/deeper."
+ Application of Lonib-Scarele analysis confirms this (obvious) period with very high confidence (e... 299.9% xienificant in the April data).," Application of Lomb-Scargle analysis confirms this (obvious) period with very high confidence (e.g., $>99.9\%$ significant in the April data)."
+ The lightcurve also shows additional reproducible structure that is seen in both March and April 2005 3.511 data: for example. there is a plateau just following the eclipse. aud. note the low-amplitude relative depression near phase 0.6.," The lightcurve also shows additional reproducible structure that is seen in both March and April 2005 3.5m data; for example, there is a plateau just following the eclipse, and note the low-amplitude relative depression near phase 0.6."
+ As fay as we are aware. SDSS 10056|3621 is the first example of an eclipsing eiAM. CVn.," As far as we are aware, SDSS J0926+3624 is the first example of an eclipsing AM CVn."
+ Our visual spectroscopic search vielded Ll new AND CVu candidates frou recent SDSS. spectroscopic plates covering about 2100 dee?. while the subsequent aleorithimic search recovered 3° AMD ονα candidates (two of the new candidates discussed herein. as wel as the Roelofs et al.," Our visual spectroscopic search yielded 4 new AM CVn candidates from recent SDSS spectroscopic plates covering about 2400 $^2$, while the subsequent algorithmic search recovered 3 AM CVn candidates (two of the new candidates discussed herein, as well as the Roelofs et al."
+ 200L object. SDSS J1210-0159) frou plates cucompassing about 1700 dee?.," 2004 object, SDSS J1240-0159) from plates encompassing about 4700 $^2$."
+ We are not aware of any examples in the SDSS spectral database of strong cluissiou-line AM. CVni binaries obviously άσσος by either visual or algorithmic searches., We are not aware of any examples in the SDSS spectral database of strong emission-line AM CVn binaries obviously missed by either visual or algorithmic searches.
+ On the other haud. the overlap-uunibers are too siiall to be considere definitive.," On the other hand, the overlap-numbers are too small to be considered definitive."
+ Accounting for the 250 plates covered both in the visual and aleorithimic searches. the combines curent SDSS) spectroscopic vield is 5 (emissiou-lue) AM CVn candidates from a region caucompassing about 5900 dee? of sy.," Accounting for the 250 plates covered both in the visual and algorithmic searches, the combined current SDSS spectroscopic yield is 5 (emission-line) AM CVn candidates from a region encompassing about 5900 $^2$ of sky."
+" We thus estimate that spectroscopically similar AAT CV systems might be expected to be found frou, the SDSS spectral database at a rough surface deusity of order 1 every 1200 dee.", We thus estimate that spectroscopically similar AM CVn systems might be expected to be found from the SDSS spectral database at a rough surface density of order 1 every 1200 $^2$.
+ Estimates of the actual surface density of SDSS AND ογι binaries (even just those with heliuu in cussion) are complicated in that the AMD CVu candidates found thus far with SDSS were chosen for spectroscopy by: 4) several different target, Estimates of the actual surface density of SDSS AM CVn binaries (even just those with helium in emission) are complicated in that the AM CVn candidates found thus far with SDSS were chosen for spectroscopy by: (i) several different target
+ | (secAcrtsrecentreviews).. (e.g.Micheletal.2008).. (Gillilandetal.2010:Chaplin2011)..," $^{-1}$ \citep[see][for recent reviews]{Aerts08, Bedding11b}. \citep[e.g.][]{Michel08}. \citep{Gilliland10, Chaplin11a}."
+ 1. /.1986): Tere. Av ds the so-called large separation between modes of the same / aud consecutive à. while ὃν is the small separation between nodes of different 7. aud € js a climensionless offset.," $n$ $l$: Here, $\Delta\nu$ is the so-called large separation between modes of the same $l$ and consecutive $n$ , while $\delta\nu_{0l}$ is the small separation between modes of different $l$, and $\epsilon$ is a dimensionless offset."
+ To a good approximation. Av is proportional to the square root of the mean density of the star (Ulrich1986). aud iu Section ?? we juvestigate the validity of this approximation.," To a good approximation, $\Delta\nu$ is proportional to the square root of the mean density of the star \citep{Ulrich86} and in Section \ref{scale} we investigate the validity of this approximation."
+ The s«1uall separations. Ongj. are scusitive to the structure of the core and hence to the age of the star. at least on the main sequence.," The small separations, $\delta\nu_{0l}$, are sensitive to the structure of the core and hence to the age of the star, at least on the main sequence."
+ These somewhat orthogonal dependencies leads to their use iu the so-called C-D diagram. in which the large aud small separations are plotted against each other (Christenscu-Dalseaarcd 1981).," These somewhat orthogonal dependencies leads to their use in the so-called C-D diagram, in which the large and small separations are plotted against each other \citep{C-D84}."
+. Calculating the C-D diagram is one of the main ais of this paper., Calculating the C-D diagram is one of the main aims of this paper.
+ Previous studies of the C-D diagram aud its variations have determined the expected evolution of stars with varving dass and metallheitv (Ulrich1986:Gough1987:Cliaisteuscu-Dalseaard 1988). and assessed the feasibility of applviug the diagram to real data (Monteroetal. 2009).," Previous studies of the C-D diagram and its variations have determined the expected evolution of stars with varying mass and metallicity \citep{Ulrich86,Gough87,C-D88}, and assessed the feasibility of applying the diagram to real data \citep{Monteiro02,OtiFloranes05,Mazumdar05,Gai09}."
+. However. none of these studies followed the evolution bevoud the end of the main sequence.," However, none of these studies followed the evolution beyond the end of the main sequence."
+ Recently. Moutalbanetal.(2010) computed the theoretical spectrum of solu-like oscillations in red-giaut stars. findine that the simall separation νο depends alinost linearly on Ar. in aerecment with the red-egiaut results fromKepler (Beddingetal.2010a:Tuber2010).," Recently, \citet{Montalban10} computed the theoretical spectrum of solar-like oscillations in red-giant stars, finding that the small separation $\delta\nu_{02}$ depends almost linearly on $\Delta\nu$, in agreement with the red-giant results from \citep{Bedding10c, Huber10}."
+. Tn Section ?? we bridee the gap. exteudiug the C-D diagram bevoud main-sequencestars to the subeiauts and uptowards the tip of the red-eiaut brauch.," In Section \ref{CD} we bridge the gap, extending the C-D diagram beyond main-sequencestars to the subgiants and uptowards the tip of the red-giant branch."
+ Acomplication with the C-D diagram for subeiauts and red-eiaut stars arises four mode bumping., Acomplication with the C-D diagram for subgiants and red-giant stars arises from mode bumping.
+ As stars, As stars
+ Here. we discuss this issue in terms of the grain size distribution.," Here, we discuss this issue in terms of the grain size distribution."
+ The contribution from grains in a logarithmic size range Bano. Ine| dine] to the extinction can be written as essi= Qx(a)n(a)a dine. where ένα). is the extinction cross section normalized to the geometrical cross section Gra7).," The contribution from grains in a logarithmic size range $\ln a$ , $\ln a +\mathrm{d}\ln a$ ] to the extinction can be written as $\mathrm{d}\kappa_\mathrm{ext}\equiv
+\pi a^2Q_\lambda (a)n(a)a\,\mathrm{d}\ln a$ , where $Q_\lambda (a)$ is the extinction cross section normalized to the geometrical cross section $\pi a^2$ )."
+" If the size distribution is approximated by a power-law (nxe.""over a certain size range. cles/dinexa? να)."," If the size distribution is approximated by a power-law $n\propto a^{-p}$ ) over a certain size range, $\mathrm{d}\kappa_\mathrm{ext}/\mathrm{d}\ln a\propto
+a^{3-p}Q_\lambda(a)$ ."
+" Since Q\(a)~1 for 2x«aA and (Qu(a)xa for σπα À te.g.Bohren&Hutf-man 1983). we obtain csi/dlnexa?"" for σπαZoA and disdinaxat"" for 2za«A."," Since $Q_\lambda (a)\sim 1$ for $2\pi a\ga\lambda$ and $Q_\lambda (a)\propto a$ for $2\pi a\ll\lambda$ \citep[e.g.][]{bohren83}, we obtain $\mathrm{d}\kappa_\mathrm{ext}/\mathrm{d}\ln a\propto a^{3-p}$ for $2\pi a\ga\lambda$ and $\mathrm{d}\kappa_\mathrm{ext}/\mathrm{d}\ln a\propto a^{4-p}$ for $2\pi a\ll\lambda$."
+ Thus. if p«3. the largest grains have the largest contribution to the extinction.," Thus, if $p<3$, the largest grains have the largest contribution to the extinction."
+ In order for small grains to have significant contribution to the extinction. pz3 should be satistied.," In order for small grains to have significant contribution to the extinction, $p\geq 3$ should be satisfied."
+ If 3.«p4. the largest contribution to the extinction comes from the grains with 27a~A.," If $3<p<4$, the largest contribution to the extinction comes from the grains with $2\pi a\sim\lambda$."
+ In other words. the UV (A~0.2 jum) extinction curve is steepened significantly if grains with e0.03μπι are produced and p3 is satisfied around this grain size.," In other words, the UV $\lambda\sim 0.2~\micron$ ) extinction curve is steepened significantly if grains with $a\sim 0.03~\micron$ are produced and $p\ga 3$ is satisfied around this grain size."
+" From reftis:size,0.1 4.. weobserecthatalargenumberofqrainswilhagnderestimate are produced and the slope around this grain radius is p<3 for the solar metallicity cases."," From \\ref{fig:size_n0.1}- \ref{fig:size_n10}, we observe that a large number of grains with $a\sim 0.03~\micron$ are produced and the slope around this grain radius is $p\ga 3$ for the solar metallicity cases."
+" Indeed. the UV slope of extinction curve is steepened for the solar metallicity cases as shown in reftig:ext,, ge.."," Indeed, the UV slope of extinction curve is steepened for the solar metallicity cases as shown in \\ref{fig:ext_age}."
+ In this paper. the shattered fragments are distributed with a size distribution with exponent a;=3.3 (Section ?2)). Jonesal. (1996)..," In this paper, the shattered fragments are distributed with a size distribution with exponent $\alpha_\mathrm{f}=3.3$ (Section \ref{subsec:shatter}) ). \citet{jones96}, ,"
+" from a discussion on the cratering flow. argue that the value of o, slightly larger than 3 is robust."," from a discussion on the cratering flow, argue that the value of $\alpha_\mathrm{f}$ slightly larger than 3 is robust."
+ Even if ay=2.5 is assumed as an extreme case. the difference in the extinction curve is less than at A=0.1pum and smaller at longer wavelengths (see the Appendix ?? for details).," Even if $\alpha_\mathrm{f}=2.5$ is assumed as an extreme case, the difference in the extinction curve is less than at $\lambda =0.1~\micron$ and smaller at longer wavelengths (see the Appendix \ref{app:alpha2.5} for details)."
+ From the results above. the presence of small grains in starburst environments is generally predicted. although SNe II tend to eject large grains because ofSNRs.," From the results above, the presence of small grains in starburst environments is generally predicted, although SNe II tend to eject large grains because of."
+ For example. BCDs tor H galaxies) in the nearby Universe host large ionized region and the age of the current star formation episode is a few Myr-20 Myr (e.g.Hirashita&Hunt2004:Takeuchial. 2005).," For example, BCDs (or H galaxies) in the nearby Universe host large ionized region and the age of the current star formation episode is a few Myr–20 Myr \citep[e.g.][]{hirashita04,takeuchi05}."
+. These ages are just in the range where shattering could modify the grain size distribution and extinction curve. although we should take into account the low metallicity in BCDs.," These ages are just in the range where shattering could modify the grain size distribution and extinction curve, although we should take into account the low metallicity in BCDs."
+ Some BCDs show an excess of near-infrared emission (e.g.Hunt.Vanzi.&Thuan 2001). which can be attributed by the emission from transiently heated very small grains CAannestad&Kenyon1979:Sellgren1984:Draine&And," Some BCDs show an excess of near-infrared emission \citep[e.g.][]{hunt01}, which can be attributed by the emission from transiently heated very small grains \citep{aannestad79,sellgren84,draine85}."
+erson 1985).. Gallianoetal.(2005) have carried out a comprehensive analysis of the SEDs of dust and stars in some dwarf galaxies (dwarf irregular galaxies and BCDs). and have shown that the grain size is biased to small grains with ~ afew nm.," \citet{galliano05} have carried out a comprehensive analysis of the SEDs of dust and stars in some dwarf galaxies (dwarf irregular galaxies and BCDs), and have shown that the grain size is biased to small grains with $\sim$ a few nm."
+ Since their sample galaxies have metallicities larger than 1/10 Z.;. shattering in WIM can work as a production source of nm-sized grains on time-scales of a few Myr and thus can be considered as an origin of small grains in these galaxies.," Since their sample galaxies have metallicities larger than 1/10 $_{\sun}$, shattering in WIM can work as a production source of nm-sized grains on time-scales of a few Myr and thus can be considered as an origin of small grains in these galaxies."
+ It is natural to expect that a similar condition tturbulence in WIM sustained more than a few Myr) is generally realized in starburst galaxies., It is natural to expect that a similar condition turbulence in WIM sustained more than a few Myr) is generally realized in starburst galaxies.
+ Although it is hard to compare the extinction curve with the observed wavelength dependence of the dust attenuation because of the effects of radiative transfer (Calzetti 2005).. shattering may be crucial to reproduce the reddening in starburst galaxies.," Although it is hard to compare the extinction curve with the observed wavelength dependence of the dust attenuation because of the effects of radiative transfer \citep{calzetti01,inoue05}, shattering may be crucial to reproduce the reddening in starburst galaxies."
+ Therefore. shattering should be considered as a source of small grains. which contribute to the reddening.," Therefore, shattering should be considered as a source of small grains, which contribute to the reddening."
+ Efficient shattering also occurs in the ISM by the passage of SN shocks., Efficient shattering also occurs in the ISM by the passage of SN shocks.
+ Jonesetal.(1996). show that a large fraction of large grains with e0.1pum is redistributed into smaller grains by a single passage of shock with a velocity of ~100 km +., \citet{jones96} show that a large fraction of large grains with $a>0.1~\micron$ is redistributed into smaller grains by a single passage of shock with a velocity of $\sim 100$ km $^{-1}$.
+ and are subject to more collisions with dust., and are subject to more collisions with dust.
+ Jonesetal.(1996). consider the MRN distribution as the initial grain size distribution. which enhances the shattering efficiency compared with our case. because of the enhanced collision with the abundant small grains.," \citet{jones96}
+ consider the MRN distribution as the initial grain size distribution, which enhances the shattering efficiency compared with our case, because of the enhanced collision with the abundant small grains."
+ Below. we estimate the time-scale on which shattering in SN shocks destroys large grains based on Jonesetal.(1996).. although the time-scale obtained might be an distributio," Below, we estimate the time-scale on which shattering in SN shocks destroys large grains based on \citet{jones96}, although the time-scale obtained might be an underestimate."
+"n, The time-scale on which shattering in SN shocks effectively destroys large grains can basically estimated by a similar way to McKee(1989).", The time-scale on which shattering in SN shocks effectively destroys large grains can basically estimated by a similar way to \citet{mckee89}.
+. A single SN can sweep Aly~107AL. of gas Mac2/2~fen with a shock velocity ος~LOO km 1 and energy given to gas by a SN y~1073 erg)., A single SN can sweep $M_\mathrm{sw}\sim 10^4~\Msun$ of gas $M_\mathrm{sw}v_\mathrm{s}^2/2\sim E_\mathrm{SN}$ with a shock velocity $v_\mathrm{s}\sim 100$ km $^{-1}$ and energy given to gas by a SN $E_\mathrm{SN}\sim 10^{51}$ erg).
+ Then. the gas muss swept by SN shocks with eZ100 km + per unit time can be estimated as AZ.7. where > is the SN rate.," Then, the gas mass swept by SN shocks with $v_\mathrm{s}\ga 100$ km $^{-1}$ per unit time can be estimated as $M_\mathrm{sw}\gamma$, where $\gamma$ is the SN rate."
+ Thus. the on which the entire gas mass Ad. is affected by shattering by SN shocks is estimated as Taw~AM(Mas).," Thus, the time-scale on which the entire gas mass $M_\mathrm{g}$ is affected by shattering by SN shocks is estimated as $\tau_\mathrm{sw}\sim M_\mathrm{g}/
+(M_\mathrm{sw}\gamma)$."
+ Since +fer~10.7M. for a Salpeter initial mass function/ (Salpeter1955) (c is the star formation rate). the above time-scale is estimated. as Teol0FAL.c.," Since $\gamma/\psi\sim 10^{-2}\Msun^{-1}$ for a Salpeter initial mass function \citep{salpeter55}
+ $\psi$ is the star formation rate), the above time-scale is estimated as $\tau_\mathrm{sw}\sim 10^{-2}M_\mathrm{g}/\psi$."
+ This estimate indicates that the shattering scale by SN shocks is about 0.01 times the gas consumption time- by star formation., This estimate indicates that the shattering time-scale by SN shocks is about 0.01 times the gas consumption time-scale by star formation.
+ In starburst environments. Ad.ευ107— 10° yr may be reasonable (Youngetal.1986)... and shattering in SN shocks occurs in 1-10 Myr. which is comparable to the time-scale investigated in this paper.," In starburst environments, $M_\mathrm{g}/\psi\sim 10^8$ $10^9$ yr may be reasonable \citep{young86}, and shattering in SN shocks occurs in 1–10 Myr, which is comparable to the time-scale investigated in this paper."
+ Therefore. both shattering in turbulence and that in SN shocks can affect the grain size distribution.," Therefore, both shattering in turbulence and that in SN shocks can affect the grain size distribution."
+ A detailed calculation of shattering in SN shocks of grains produced by SNe IL is required before we judge which of these two shattering mechanisms 1s dominated., A detailed calculation of shattering in SN shocks of grains produced by SNe II is required before we judge which of these two shattering mechanisms is dominated.
+ It might be also useful to discuss our results in terms of the extinction curves of the Large and Small Magellanic Cloud (LMC and SMC). both of which have developed H regions such as 30 Doradus.," It might be also useful to discuss our results in terms of the extinction curves of the Large and Small Magellanic Cloud (LMC and SMC), both of which have developed H regions such as 30 Doradus."
+ Indeed. Bernardetal.(2008) indicate that the 70 jum excess around 30 Doradus can be explained by an enhancement of the abundance of very small grains possibly by the destruction of large grains.," Indeed, \citet{bernard08} indicate that the 70 $\micron$ excess around 30 Doradus can be explained by an enhancement of the abundance of very small grains possibly by the destruction of large grains."
+ Botetal.(2004). find this excess in the SMC., \citet{bot04} find this excess in the SMC.
+ Paradisetal.(2009). show that the very small grain abundance is really enhanced around 30 Doradus by using an SED model of dust emission., \citet{paradis09} show that the very small grain abundance is really enhanced around 30 Doradus by using an SED model of dust emission.
+ However. the extinction curves in these galaxies are much steeperthan our results Glfelyc2.9 and 3.2 at A— for the LMC and the SMC. respectively: Pei 19923).," However, the extinction curves in these galaxies are much steeperthan our results $A_\lambda /A_V\simeq 2.9$ and 3.2 at $\lambda\simeq 0.2~\micron$ for the LMC and the SMC, respectively; \citealt{pei92}) )."
+ Since those galaxies have less intensestar formation than BCDs. it is hard to extract the starbursting components where shattering of large grains should be working as investigated in this paper.," Since those galaxies have less intensestar formation than BCDs, it is hard to extract the starbursting components where shattering of large grains should be working as investigated in this paper."
+ The steep extinction curves of the LMC and the SMC indicate that we should consider not only the dust production/shattering in star-forming regions but also some other mechanisms which act as efficient production sources of small grains., The steep extinction curves of the LMC and the SMC indicate that we should consider not only the dust production/shattering in star-forming regions but also some other mechanisms which act as efficient production sources of small grains.
+ For example. shattering in," For example, shattering in"
+2010).,.
+.. As flix is being redistributed owing to the granular evolution. the bundles are dispersed and the spatial smearing of more isolated loop-like structures reduces the linear polarization signal to values below the noise level.," As flux is being redistributed owing to the granular evolution, the bundles are dispersed and the spatial smearing of more isolated loop-like structures reduces the linear polarization signal to values below the noise level."
+ Based on Sunrise/IMaX data and using an automated detection method. we obtained statistical properties of 4536 features with significant linear polarization signal.," Based on Sunrise/IMaX data and using an automated detection method, we obtained statistical properties of 4536 features with significant linear polarization signal."
+ Their iletimes are consistent with examples given previously in the literature., Their lifetimes are consistent with examples given previously in the literature.
+ However. the iletime distribution indicates no characteristic value. in contrast (o. previous studies (Ishikawaetal.2008:Ishikawa&Tsuneta2009:Jin2009).," However, the lifetime distribution indicates no characteristic value, in contrast to previous studies \citep{Ishikawa:etal:2008,Ishikawa:Tsuneta:2009,Jin:etal:2009}."
+. The detected. features iive no characteristic size either., The detected features have no characteristic size either.
+ Around 97% of them are smaller than 1 arcesec?. which is (he value previously (taken as (he mean size of HIF (shikawa&Tsuneta2009).," Around $97$ of them are smaller than $\sim$ 1 $^2$, which is the value previously taken as the mean size of HIF \citep{Ishikawa:Tsuneta:2009}."
+. We lind (hat their rate of occurrence is 1-2 orders of magnitude higher (han reported earlier (Litesetal.1996:Ishikawa&Tsuneta2009;MartínezGonzálezDellotRaibio2009).," We find that their rate of occurrence is 1-2 orders of magnitude higher than reported earlier \citep{Lites:etal:1996,Ishikawa:Tsuneta:2009,Marian:Luis:2009}."
+. We attribute Chis discrepancy (o selection effects., We attribute this discrepancy to selection effects.
+" If we take only the biggest features (sizes >0.55 arcsec?). only 45 of the detected features remain and the rate of occurrence decreases to ~4-10? | 7, which is in closer agreement with the references cited above."," If we take only the biggest features (sizes $>0.88$ $^2$ ), only $\sim$ of the detected features remain and the rate of occurrence decreases to $\sim4\cdot10^{-5}$ $^{-1}$ $^{-2}$, which is in closer agreement with the references cited above."
+ Longer-lived HIF tend to be lareer and display a higher mean linear polarization signals., Longer-lived HIF tend to be larger and display a higher mean linear polarization signals.
+ The HIF appear preferentially al (he eranule boundaries. with most of them being caught by downllows at some point in their evolution.," The HIF appear preferentially at the granule boundaries, with most of them being caught by downflows at some point in their evolution."
+ We showed that ~16% of the features we detected are completely embedded in upflows and ~8% are entirely embedded in downllows., We showed that $\sim$ of the features we detected are completely embedded in upflows and $\sim$ are entirely embedded in downflows.
+ The latter are very. small in size (as illustrated by the two examples discussed in greater detail)., The latter are very small in size (as illustrated by the two examples discussed in greater detail).
+ Although their origin is still uncertain it is clear that they do not fit into (he scenario of magnetic [Iux emergence as their plvsical cause., Although their origin is still uncertain it is clear that they do not fit into the scenario of magnetic flux emergence as their physical cause.
+"comes from the ""C destruction via the CN cycle, and the oxygen content O remains about constant.","comes from the $^{12}$ C destruction via the CN cycle, and the oxygen content $O$ remains about constant."
+" Thus, one has dC=-$dN since N globally results from the addition of two protons to !7C, This gives the slope always in mass fractions."," Thus, one has ${\rm d}C= -\frac{6}{7} {\rm d}N $ since $^{14}$ N globally results from the addition of two protons to $^{12}$ C, This gives the slope always in mass fractions."
+" With ratios N/C and N/O of 0.31 and 0.11, respectively, we get This ratio is evidently greater than 1, since as N starts growing, C decreases, while O does not vary much."," With ratios $N/C$ and $N/O$ of 0.31 and 0.11, respectively, we get This ratio is evidently greater than 1, since as $N$ starts growing, $C$ decreases, while $O$ does not vary much."
+ The relation turns slightly upward as N/C is increasing owing to the term in brackets in Eq. (2))., The relation turns slightly upward as $N/C$ is increasing owing to the term in brackets in Eq. \ref{slope}) ).
+" However, at some advanced stage in evolution, corresponding to WN stars not shown here, the curve will saturate and turn down slightly (Maeder2009,p. 699),, since the CN cycle is then at equilibrium, while !6O is still turned to '4N. Dilution mixes a fraction f of N++AN enriched and C depleted materials with a fraction (1—f) of the original N and C."," However, at some advanced stage in evolution, corresponding to WN stars not shown here, the curve will saturate and turn down slightly \citep[p.~699]{maeder09}, , since the CN cycle is then at equilibrium, while $^{16}$ O is still turned to $^{14}$ N. Dilution mixes a fraction $f$ of $N$ $\Delta N$ enriched and C depleted materials with a fraction $(1-f)$ of the original $N$ and $C$."
+" Under the same assumptions as above, it is easy to show that, to the first order, the slope for the relative enrichments in the N/C vs. N/O plot behaves the same way as in Eq. (2))"," Under the same assumptions as above, it is easy to show that, to the first order, the slope for the relative enrichments in the $N/C$ vs. $N/O$ plot behaves the same way as in Eq. \ref{slope}) )"
+ independently of f., independently of $f$.
+" The value of f, however, determines the amplitudes of the departures from the cosmic ratios."," The value of $f$, however, determines the amplitudes of the departures from the cosmic ratios."
+" Our models with rotational mixing (Meynet Maeder 2003, MM03; Ekstrómm et al."," Our models with rotational mixing (Meynet Maeder 2003, MM03; Ekströmm et al."
+" 2008, E08) or with rotation and magnetic fields (MMOS), as illustrated e.g. in Fig. 1,,"," 2008, E08) or with rotation and magnetic fields (MM05), as illustrated e.g. in Fig. \ref{litsummary},"
+" have an initial slope EI 944, which is in excellent agreement with Eq. (3))."," have an initial slope $\frac{{\rm d}(N/C)}{{\rm
+d}(N/O)}$ $\approx$ 4, which is in excellent agreement with Eq. \ref{sl4}) )."
+" The amplitude f of the mixing depends on the various model assumptions, in particular on the treatment of the shear mixing with or without horizontal turbulence."," The amplitude $f$ of the mixing depends on the various model assumptions, in particular on the treatment of the shear mixing with or without horizontal turbulence."
+" The models without horizontal turbulence (Meynet&Maeder2000,MMO0) predict more mixing than models that account for it (MM03)."," The models without horizontal turbulence \citep[MM00]{mema00}
+ predict more mixing than models that account for it (MM03)."
+ Models that include both rotation and magnetic field predict a still larger mixing (ΜΜΟΣ)., Models that include both rotation and magnetic field predict a still larger mixing (MM05).
+" Let us now consider the behaviour of the helium surface content Y, vs. N/O (as illustrated later).", Let us now consider the behaviour of the helium surface content $Y_{\mathrm{s}}$ vs. $N/O$ (as illustrated later).
+" Strictly and only at the very beginning of the CN burning, and under the assumption of an initially constant oxygen, we get dY, = 2dN, since when 4 units of mass of helium are made, 14 units of mass of nitrogen are produced."," Strictly and only at the very beginning of the CN burning, and under the assumption of an initially constant oxygen, we get ${\rm d}Y_{\mathrm{s}}$ $=$ $\frac{2}{7}{\rm d}N$, since when 4 units of mass of helium are made, 14 units of mass of nitrogen are produced."
+" The slope is i.e., it 7is essentially flat initially."," The slope is i.e., it is essentially flat initially."
+" Later in the evolution, both N and O change simultaneously, and one has to rely on numerical models."," Later in the evolution, both N and O change simultaneously, and one has to rely on numerical models."
+" The resulting slope in the models can vary; e.g., a steeper slope is inferred for the Μο model of both MMO3 and models byMMO0 than for the Μο models of MMO03 and MMOS."," The resulting slope in the models can vary; e.g., a steeper slope is inferred for the $M_{\odot}$ model of both MM03 and models byMM00 than for the $M_{\odot}$ models of MM03 and MM05."
+ This depends on whether the matter that arrives at the surface comes from inner regions that are at both, This depends on whether the matter that arrives at the surface comes from inner regions that are at both
+data revealed that the source is a pulsar withs.,data revealed that the source is a pulsar with.
+. The spectral fit provided a value of Ny=3.2x10? cm? (2)., The spectral fit provided a value of $N_{H} = 3.2 \times 10^{22}$ $^{-2}$ \citep{bamba03}.
+". A fast outburst observed by INTEGRAL was attributed by ? to this A Chandra observation of the field revealed the counterpart to be 2MASS 18410043-0535465, a reddened star with weak Ha in emission (?), suggesting it was a Be star. ?,,"," A fast outburst observed by INTEGRAL was attributed by \citet{halgot04} to this A Chandra observation of the field revealed the counterpart to be 2MASS 18410043-0535465, a reddened star with weak $\alpha$ in emission \citep{halpern04}, suggesting it was a Be star. \citet{negue06},"
+" from optical spectroscopy, proposed the star is instead a luminous BO-1 type, although with some uncertainty, classifying the system as an Figure 4 shows the K, spectrum we obtained, with identified spectral features marked."," from optical spectroscopy, proposed the star is instead a luminous B0-1 type, although with some uncertainty, classifying the system as an Figure \ref{fig:AXJ1841} shows the $K_{s}$ spectrum we obtained, with identified spectral features marked."
+" The spectrum shows He1205581 eemission, accompanied by a spurious feature, possibly due to poor telluric component removal; we observe absorption atHer21 1126, a weak Nm (Cur) 211155 eemission; moreover, there is moderately strong Bry absorption."," The spectrum shows 581 emission, accompanied by a spurious feature, possibly due to poor telluric component removal; we observe absorption at 126, a weak ) 155 emission; moreover, there is moderately strong $\gamma$ absorption."
+" The side features of the Bry absorption profile are probably due to poor telluric correction, but they do not prevent us from measuring the equivalent The observed transitions are typical of an early supergiant, and by comparison with the atlases from ??,, we can conclude that the star is of B1 Ib type."," The side features of the $\gamma$ absorption profile are probably due to poor telluric correction, but they do not prevent us from measuring the equivalent The observed transitions are typical of an early supergiant, and by comparison with the atlases from \citet{hanson96,hanson05}, we can conclude that the star is of B1 Ib type."
+" Together with X-ray properties, this NIR spectral classification allows us to confirm the nature of the system as an SEXT."," Together with X-ray properties, this NIR spectral classification allows us to confirm the nature of the system as an SFXT."
+" The wind-accreting system 4U 1907409 (?) is a known HMXB consisting of a neutron star in an eccentric (e = 0.28) 8.3753 dayorbit around its companion, which has been optically identified as a highly reddened star (?).."," The wind-accreting system 4U 1907+09 \citep{giacc71} is a known HMXB consisting of a neutron star in an eccentric $e$ = 0.28) 8.3753 dayorbit around its companion, which has been optically identified as a highly reddened star \citep{schwa80}."
+ The spectral classification of the counterpart to 4U 1907+09 has been matter of debate., The spectral classification of the counterpart to 4U 1907+09 has been matter of debate.
+ The presence of X-ray flaring seen twice per neutron star orbit (?) had led some authors (e.g.???) to the hypothesis of a Be star companion.," The presence of X-ray flaring seen twice per neutron star orbit \citep{mar80} had led some authors \citep[e.g.][]{maki84,copa87,Iye86} to the hypothesis of a Be star companion."
+" However, this classification would require a distance of <1.5 kpc, which is in contradiction to the significant interstellar extinction measured in optical observations by ?,, who also classified the counterpart as a B supergiant."," However, this classification would require a distance of $<$ 1.5 kpc, which is in contradiction to the significant interstellar extinction measured in optical observations by \citet{vanker89}, who also classified the counterpart as a B supergiant."
+" Using interstellar atomic lines of Na I and K I, set a lower limit of 5 kpc for the distance and proposed that the stellar companion is instead a O8-O9 Ia supergiant with an effective temperature of 305500 K, a radius of 26 Ro, a luminosity of 5x10? Lo, and a mass loss rate of 7x10-° Mo ντ,"," Using interstellar atomic lines of Na I and K I, \citet{cox05} set a lower limit of 5 kpc for the distance and proposed that the stellar companion is instead a O8-O9 Ia supergiant with an effective temperature of 500 K, a radius of 26 $R_{\sun}$, a luminosity of $5 \times 10^5$ $_{\sun}$, and a mass loss rate of $7 \times 10^{-6}$ $_{\sun}$ $^{-1}$."
+" Similarly to other accreting neutron stars, the X-ray continuum of 4U 1907+09 can be described by a power-law spectrum with an exponential turnover at 13 keV. The spectrum is modified by strong photoelectric absorption with a column density Ny=1.5—5.7x10? cm""? (e.g.?).."," Similarly to other accreting neutron stars, the X-ray continuum of 4U 1907+09 can be described by a power-law spectrum with an exponential turnover at 13 keV. The spectrum is modified by strong photoelectric absorption with a column density $N_{H} = 1.5 - 5.7 \times 10^{22}$ $^{-2}$ \citep[e.g.][]{copa87}."
+" We show for the first time an infrared spectrum of the source, which permits us to confirm the spectral classification as estimated from optical data."," We show for the first time an infrared spectrum of the source, which permits us to confirm the spectral classification as estimated from optical data."
+" Figure 5 presents the K, spectrum we obtained, with identified spectral features marked."," Figure \ref{fig:4U1907} presents the $K_{s}$ spectrum we obtained, with identified spectral features marked."
+" The spectrum shows absorption both at 205580 aand at 1126A,, a weak N ΠΙ (or C III) emission line at 211155 aand strong Bry absorption (EW « 4 A)), the typical features of an early supergiant."," The spectrum shows absorption both at 580 and at 126, a weak N III (or C III) emission line at 155 and strong $\gamma$ absorption (EW $<$ 4 ), the typical features of an early supergiant."
+ The presence of1 5580 in absorption strongly constrains the spectral type to a late O star., The presence of 580 in absorption strongly constrains the spectral type to a late O star.
+" By comparison with the atlases from ??,, we conclude that the star isan O9.5 Iab."," By comparison with the atlases from \citet{hanson96,hanson05}, we conclude that the star isan O9.5 Iab."
+ We thus confirm and refine the previous spectral classification., We thus confirm and refine the previous spectral classification.
+ The INTEGRAL discovery of this source was reported by ?.., The INTEGRAL discovery of this source was reported by \citet{hann03}. .
+ Observations with the Rossi X-Ray Timing Explorer (RXTE), Observations with the Rossi X-Ray Timing Explorer (RXTE)
+We also find a small number of opticallv-faint. very low redshift. compact objects which fall outside the genera trend in the by> plane.,"We also find a small number of optically-faint, very low redshift, compact objects which fall outside the general trend in the $b_{\rm J}-z$ plane."
+ The X-ray luminositics of these sources range from 4107 cerg/s to 1107 cores suggesting that they are either. very luminous starbursts or AGN., The X-ray luminosities of these sources range from $4\times 10^{41}$ erg/s to $1\times 10^{43}$ erg/s suggesting that they are either very luminous starbursts or AGN.
+ However. in all but one case. the 6dEGS optica spectra does not provide any evidence of AGN or starburs activity.," However, in all but one case, the 6dFGS optical spectra does not provide any evidence of AGN or starburst activity."
+ Further observational follow up is needed to confirm the physical properties of these sources., Further observational follow up is needed to confirm the physical properties of these sources.
+ Vhere are 918. (27%) RASS6dEGS. sources. detecte in either the llz NRAO VLA Sky Survey (NVSS) or the MMlIz Sydney. University Molonglo Sky Survey (SUMSS)., There are 918 $\%$ ) RASS–6dFGS sources detected in either the GHz NRAO VLA Sky Survey (NVSS) or the MHz Sydney University Molonglo Sky Survey (SUMSS).
+ The fraction of sources with radio counterparts changes with redshift. and at z21 nearly all the RASS6dECGS sources have radio detections.," The fraction of sources with radio counterparts changes with redshift, and at $z>1$ nearly all the RASS–6dFGS sources have radio detections."
+ These sources are strong radio sources with a median [lux density. of 1151muns whereas the median [lux censity for the ful sample is mmy., These sources are strong radio sources with a median flux density of mJy whereas the median flux density for the full sample is mJy.
+ We attribute this to the presence of a radio jet which Doppler boosts the radio emission., We attribute this to the presence of a radio jet which Doppler boosts the radio emission.
+ The ray [lux of these sources is also boosted by a jet componen and thus at large redshifts selecting bright N-rav. sources preferentially selects radio-oud ACGN., The X-ray flux of these sources is also boosted by a jet component and thus at large redshifts selecting bright X-ray sources preferentially selects radio-loud AGN.
+ The RASSGdb€GS catalogue. when reduced. to jus those sources with radio detections. can be used as the southern counterpart to the RBSCNVSS catalogue and as such ollers a large sample of BL-Lac anc blazar sources.," The RASS–6dFGS catalogue, when reduced to just those sources with radio detections, can be used as the southern counterpart to the RBSC–NVSS catalogue and as such offers a large sample of BL-Lac and blazar sources."
+ Other. properties. of RASSος sources. in. particular the high-frequeney. radio properties. will be examined. in forthcoming papers.," Other properties of RASS–6dFGS sources, in particular the high-frequency radio properties, will be examined in forthcoming papers."
+ This will enable further. studies: of BL-Lac objects and a more extensive. analysis of the multiwavelength properties (X-ray - optical - radio)., This will enable further studies of BL-Lac objects and a more extensive analysis of the multiwavelength properties (X-ray - optical - radio).
+ Future work also includes observational follow-up of the small number of optically faint. very low redshift sources identified in this paper.," Future work also includes observational follow-up of the small number of optically faint, very low redshift sources identified in this paper."
+ 2 , \ref{wrongredshiftstab} 
+"is calculated by where dy, is the luminosity distance and S, the fluence in the1/(1+z) keV to 10/(1+z) MeV frame.",is calculated by where $d_L$ is the luminosity distance and $S_\gamma$ the fluence in the$1/(1+z)$ keV to $10/(1+z)$ MeV frame.
+" We determine S, using the energy flux provided by the best-fit spectral parameters and multiplying it with the total time interval over which the fit was performed.", We determine $S_\gamma$ using the energy flux provided by the best-fit spectral parameters and multiplying it with the total time interval over which the fit was performed.
+" Since we performed the fit for time intervals where the count rate exceeded a S/N ratio of 3.5, it happened that some time intervals of some bursts were not included in the fit (e.g. phases of quiescence where the count rate dropped back to the background level)."," Since we performed the fit for time intervals where the count rate exceeded a S/N ratio of 3.5, it happened that some time intervals of some bursts were not included in the fit (e.g. phases of quiescence where the count rate dropped back to the background level)."
+ These time intervals were not used to calculate the fluence., These time intervals were not used to calculate the fluence.
+" The S, distribution is shown in Fig.5..", The $S_\gamma$ distribution is shown in \ref{fig:flugbm}.
+ The median value of the fluence distribution is 1.6x10? erg cm? and the mean value is 5.9x10? erg cm?., The median value of the fluence distribution is $1.6\times10^{-5}$ erg $^{-2}$ and the mean value is $5.9\times10^{-5}$ erg $^{-2}$.
+ A log-normal fit to the data peaks at 2.2x102 erg cm?., A log-normal fit to the data peaks at $2.2\times10^{-5}$ erg $^{-2}$.
+ The distribution is shown in Fig.6.., The distribution is shown in \ref{fig:histeiso}. .
+" The distribution for the long bursts has a median and mean value of 1.2x erg and 1.4x10? erg, respectively."," The distribution for the long bursts has a median and mean value of $1.2\times10^{53}$ erg and $1.4\times10^{53}$ erg, respectively."
+" Short bursts, on the other hand, have significantly lower values of 2.9x10?! erg and 4.0x10?! erg, respectively."," Short bursts, on the other hand, have significantly lower values of $2.9\times10^{51}$ erg and $4.0\times10^{51}$ erg, respectively."
+ A log-normal fit to the long bursts reveals a central value of 105?! erg., A log-normal fit to the long bursts reveals a central value of $10^{53.1}$ erg.
+ Because our sample is dominated by long GRBs a log-normal fit to the whole distribution results in an essentially unchanged peak value (1053 erg)., Because our sample is dominated by long GRBs a log-normal fit to the whole distribution results in an essentially unchanged peak value $10^{53}$ erg).
+ ? first showed that there is a tight correlation between and (the isotropic equivalent bolometric energy determined in the energy range between 1 keV to 10 MeV)., \citet{amati02} first showed that there is a tight correlation between and (the isotropic equivalent bolometric energy determined in the energy range between 1 keV to 10 MeV).
+" This relation is now known as the ""Amati relation"".", This relation is now known as the “Amati relation”.
+ In Fig.7 we show the Amati relation for the 30 GBM GRBs with measured andΕικο., In \ref{fig:amati} we show the Amati relation for the 30 GBM GRBs with measured and.
+. While there is an evident correlation between these two quantities (Spearman's rank correlation of p=0.74 with a chance probability of 1.7x 107?) the extrinsic scatter of the long GRBs is larger by a factor of ~2 in log-space compared to ?.., While there is an evident correlation between these two quantities (Spearman's rank correlation of $\rho= 0.74$ with a chance probability of $1.7\times10^{-5}$ ) the extrinsic scatter of the long GRBs is larger by a factor of $\sim 2$ in $\log$ -space compared to \citet{amati10}.
+" Also, the best fit to our data is shifted to slightly larger Values."," Also, the best fit to our data is shifted to slightly larger values."
+ The best fit power-law index to the long GRBs of our sample is 0.52+0.06 which is in agreement with the indices obtained by e.g. ???..," The best fit power-law index to the long GRBs of our sample is $0.52\pm0.06$ which is in agreement with the indices obtained by e.g. \citet{amati10, ghina09, ghirlanda10}."
+" As has been shown by other authors in the past (seee.g.???) short bursts do not follow the relation, being situated well outside the 2 o scatter around the best-fit."," As has been shown by other authors in the past \citep[see e.g.][]{amati08, ghina09, amati10} short bursts do not follow the relation, being situated well outside the 2 $\sigma$ scatter around the best-fit."
+" This is true also for the power-law fit derived here (see Fig.7)) except for GRB 100816A. However, as already stated above this burst may actually fall in an intermediate or hybrid class of short GRBs with extended emission (seee.g. ??).."," This is true also for the power-law fit derived here (see \ref{fig:amati}) ) except for GRB 100816A. However, as already stated above this burst may actually fall in an intermediate or hybrid class of short GRBs with extended emission \citep[see e.g.][]{norris06, zhang09}. ."
+" ? found a tight correlation between the rest frame peak energy in the vF, spectrum and the 1-s peak luminosity (119) in GRBs (so called Yonetoku relation).", \citet{yonetoku04} found a tight correlation between the rest frame peak energy in the $\nu \rm{F}_{\nu}$ spectrum and the 1-s peak luminosity $(L_p)$ in GRBs (so called Yonetoku relation).
+ The peak luminosity is calculatedwith, The peak luminosity is calculatedwith
+sources with radio detections and those without.,sources with radio detections and those without.
+ What is immecdiately obvious in these plots is that at high recshilts virtually all of the sources fall into the radio sample and in [act ave quite strong radiosources!., What is immediately obvious in these plots is that at high redshifts virtually all of the sources fall into the radio sample and in fact are quite strong radio.
+. Phe median Dux density for sources with z lis H151mnis while the median Hux density for the entire RASS6dECGS radio sample is mum., The median flux density for sources with $z>1$ is mJy while the median flux density for the entire RASS–6dFGS radio sample is mJy.
+v. This suggests that the more distant sources are Doppler boosted due to the presence of a radio jet. pointed towards our line of sight., This suggests that the more distant sources are Doppler boosted due to the presence of a radio jet pointed towards our line of sight.
+ As such. the X-ray. emission (or some fraction of it) is also being boosted by a jet component and hence we only detect bright X-ray sources that are raclio-loud at high redshifts.," As such, the X-ray emission (or some fraction of it) is also being boosted by a jet component and hence we only detect bright X-ray sources that are radio-loud at high redshifts."
+ This is in agreement with ? who founcl that radio-loud QSOs have a higher average X-ray luminosity., This is in agreement with \citet{1987ApJ...313..596W} who found that radio-loud QSOs have a higher average X-ray luminosity.
+ The fraction of sources with radio detections changes with redshift. as shown in Figure 11.," The fraction of sources with radio detections changes with redshift, as shown in Figure \ref{radiodetz}."
+ Vhere is a high detection. rate at. low redshift where we detect a large fraction of low luminosity. radio-quiet. AGN in addition to the racio-Ioud sources.," There is a high detection rate at low redshift where we detect a large fraction of low luminosity, radio-quiet AGN in addition to the radio-loud sources."
+ This drops olf quite rapidly as these racio-quict sources fall below the detection limit of the radio surveys. leaving only the radio-Ioud AGN in the sample.," This drops off quite rapidly as these radio-quiet sources fall below the detection limit of the radio surveys, leaving only the radio-loud AGN in the sample."
+The initial condition for the magnetic field in this case is a two-dimensional potential field which is independent of the ν- (see Figs. 2..3.,"The initial condition for the magnetic field in this case is a two-dimensional potential field which is independent of the $x$ -coordinate (see Figs. \ref{fig:coords},"
+ for the coordinate system) and falls off exponentially with height v., \ref{fig:emergeschematic} for the coordinate system) and falls off exponentially with height $y$.
+ It is generated by the vector potential where ;Ur is the width of the arcades and οἱ is the scale height of B=, It is generated by the vector potential where $\lambda\pi$ is the width of the arcades and $\lambda$ is the scale height of $B=\abs{\vec{B}}$.
+" By is the field strength at the intersection of the lower boundaryπι +=5, with the central axis (vy).", $B_\bnd$ is the field strength at the intersection of the lower boundary $r=\rb$ with the central axis $y$ ).
+ To avoid numerical problems at thelateral boundaries. we limit the number of ares to 4. setting A=B=0 for [5]>Zurand for |]>2;Ur.," To avoid numerical problems at thelateral boundaries, we limit the number of arcs to 4, setting $\vec{A}=\vec{B}=0$ for $\abs{z} > 2\lambda\pi$and for $\abs{x} >
+2\lambda\pi$."
+ The magnetic field is embedded in aspherically stratified atmosphere with p«ση and po77. held in hydrostatic equilibrium by the staticgravitational field ᾧo+7! of a point mass at the coordinate origin.," The magnetic field is embedded in aspherically stratified atmosphere with $p
+\propto r^{-4}$ and $\rho \propto r^{-3}$, held in hydrostatic equilibrium by the staticgravitational field $\Phi \propto r^{-1}$ of a point mass at the coordinate origin."
+" Temperature and sound speed vary às Toxr and e,&777, respectively."," Temperature and sound speed vary as $T \propto r^{-1}$ and $\cs \propto r^{-1/2}$, respectively."
+" In the outer. unmagnetized region ([x[.]z|> 2,Ur). we compensated for the absence of magnetic pressure by increasing the gas pressure."," In the outer, unmagnetized region $\abs{x},\abs{z} > 2\lambda\pi$ ), we compensated for the absence of magnetic pressure by increasing the gas pressure."
+ To maintain hydrostatic equilibrium. the density is also increased correspondingly.," To maintain hydrostatic equilibrium, the density is also increased correspondingly."
+" At the lower(r2 ry) boundary we maintain. through ""ghost cells” outside of the computational domain. an azimuthal velocity field i=Vee, corresponding to rigid rotation: v—ViNSR/Ry for R€Ry and 0 elsewhere."," At the lower$r=\rb$ ) boundary we maintain, through “ghost cells” outside of the computational domain, an azimuthal velocity field $\vec{v}=v_\varphi
+\evarphi$ corresponding to rigid rotation: $v_\varphi = \vphimax R/\Rb$ for $R
+\le \Rb$ and 0 elsewhere."
+ All quantities except for B are fixed at their initial values in the ghost cells: B is extrapolated from the interior of the domain., All quantities except for $\vec{B}$ are fixed at their initial values in the ghost cells; $\vec{B}$ is extrapolated from the interior of the domain.
+ At the sides (0 and 6) and top (upper r) of the domain.we use open boundary conditions which allow for an almost force-free outflow of material and cause no evident artifacts in the form of reflections.," At the sides $\theta$ and $\phi$ ) and top (upper $r$ ) of the domain,we use open boundary conditions which allow for an almost force-free outflow of material and cause no evident artifacts in the form of reflections."
+" A 18 chosen such that in the x=0 plane B, is positive for ||<Rp and negative for Ry«[:]<285. 1.8. arcades which start at |z|Αι from the r=ην surface have their second footpoint outside the rotating region."," $\lambda$ is chosen such that in the $x=0$ plane $B_y$ is positive for $\abs{z}
+< \Rb$ and negative for $\Rb < \abs{z} < 2\Rb$, i.e. arcades which start at $\abs{z} < \Rb$ from the $r=\rb$ surface have their second footpoint outside the rotating region."
+ This is achieved by choosing the length scale of the arcade as 4|22Rj/7., This is achieved by choosing the length scale of the arcade as $\lambda = 2\Rb / \pi$.
+ In this case. the initial condition is again an equilibrium stratification with p«r7. p«r? and «77.," In this case, the initial condition is again an equilibrium stratification with $p \propto r^{-4}$, $\rho \propto r^{-3}$ and $\Phi \propto r^{-1}$."
+ However. unlike in setup D. the atmosphere is completely unmagnetized.," However, unlike in setup D, the atmosphere is completely unmagnetized."
+ The magnetic field enters the domain through the lower boundary r=n the conditions of which are determined through ghost cells at r€ry.," The magnetic field enters the domain through the lower boundary $r=\rb$, the conditions of which are determined through ghost cells at $r<\rb$."
+" There. we impose a transverse magnetic field By (= B.) with constant amplitude B, and a polarity that alternates with =(ycos4 in step sizes of 2R,. the diameter of the rotating surface (described below)."," There, we impose a transverse magnetic field $B_\theta$ $\approx B_z$ ) with constant amplitude $B_\bnd$ and a polarity that alternates with $z=\rb \cos\theta$ in step sizes of $2\Rb$, the diameter of the rotating surface (described below)."
+" Where the polarity of By changes. the solenoidality of the magnetic field is maintained by an appropriate B,."," Where the polarity of $B_\theta$ changes, the solenoidality of the magnetic field is maintained by an appropriate $B_r$."
+ The field lines thus have the shape sketched in Fig. 3.., The field lines thus have the shape sketched in Fig. \ref{fig:emergeschematic}.
+ An approximate equilibrium is maintained through lowering the gas pressure by the value of the magnetic pressure. as far as this ts possible (p must not be negative).," An approximate equilibrium is maintained through lowering the gas pressure by the value of the magnetic pressure, as far as this is possible $p$ must not be negative)."
+ For this to work. Bp=Sxpp/B; must be greater than one. which limits the possible field strengths to By<1 in our system of units (described below).," For this to work, $\betab \coloneqq 8\pi p_\bnd / B_\bnd^2$ must be greater than one, which limits the possible field strengths to $B_\bnd < 1$ in our system of units (described below)."
+ To model the emergence of magnetic fields into the atmosphere we impose the radial velocity field in the ghostcells of the lower boundary., To model the emergence of magnetic fields into the atmosphere we impose the radial velocity field in the ghostcells of the lower boundary.
+" The injection velocity is too small to form a jet by itself: the maximum amounts to 41% of the escape velocity at the boundary. gas with this speed gets theoretically as far as r=1.27, without acceleration other than gravity."," The injection velocity is too small to form a jet by itself: the maximum amounts to $41\%$ of the escape velocity at the boundary, gas with this speed gets theoretically as far as $r \approx
+1.2 \rb$ without acceleration other than gravity."
+ In addition to the radial velocity field. we maintain an azimuthal velocity field v.οR (rigid rotation) within R< Ry.," In addition to the radial velocity field, we maintain an azimuthal velocity field $v_\varphi \propto R$ (rigid rotation) within $R\le\Rb$ ."
+ The shape of the emerging magnetic field is sketched in Fig., The shape of the emerging magnetic field is sketched in Fig.
+ laa. For the other boundaries. we use the same outflow conditions as in the model described in the preceding section.," \ref{fig:emergecases}a a. For the other boundaries, we use the same outflow conditions as in the model described in the preceding section."
+ Two variations of the above-described setup have also been studied. correspondi£ to the cases (b) and (c) deseribed in Sect. 2..," Two variations of the above-described setup have also been studied, corresponding to the cases (b) and (c) described in Sect. \ref{sec:models}."
+" In the first. the positions where B, changes its polarity. and with it the radial velocity field that injects the magnetic field. are shifted by Ry in —z direction."," In the first, the positions where $B_\theta$ changes its polarity, and with it the radial velocity field that injects the magnetic field, are shifted by $R_\bnd$ in $-z$ direction."
+ Thus. both footpoints of an individual field loop rotate with the same angular velocity.," Thus, both footpoints of an individual field loop rotate with the same angular velocity."
+" In the second case (c). the magnetic field and vertical velocity in the lower boundary are changed such that only a small. confined arcade in 0.2«R/R,0.8 with awidth of 0.2R, emerges. and the azimuthal velocity field has a Keplerian profile with v...=v""VO.IRS/R for 0.1<R/Ry1."," In the second case (c), the magnetic field and vertical velocity in the lower boundary are changed such that only a small, confined arcade in $0.2 < R/R_\bnd < 0.8$ with awidth of $0.2\Rb$ emerges, and the azimuthal velocity field has a Keplerian profile with $v_\varphi = \vphimax
+\sqrt{0.1\Rb/R}$ for $0.1 \le R/R_\bnd \le 1$."
+" The inner edge (R= 0.2R,) of the arcade rotates twice as fast as the outer edge (R= 0.8A,). the difference in rotation velocity iseVB."," The inner edge $R=0.2\Rb$ ) of the arcade rotates twice as fast as the outer edge $R=0.8\Rb$ ), the difference in rotation velocity is $\vphimax/\sqrt{8}$."
+" The models described above contain the following 6 parameters. not all of which are independent: ry. Rp. pp. py. By and v""."," The models described above contain the following 6 parameters, not all of which are independent: $r_\bnd$ , $\Rb$ , $\rho_\bnd$ , $p_\bnd$ , $B_\bnd$ and $\vphimax$ ."
+ We eliminate the dependences by taking /;= 28s. Po=py and po=py as units of length. density and pressure. and expressing all physical quantities in terms of these. see Table 1..," We eliminate the dependences by taking $l_0 \equiv 2\Rb$ $\rho_0
+\equiv \rho_\bnd$ and $p_0 \equiv p_\bnd$ as units of length, density and pressure, and expressing all physical quantities in terms of these, see Table \ref{tab:units}. ."
+ The remaining3 parameters can then be expressed as dimensionless numbers., The remaining3 parameters can then be expressed as dimensionless numbers.
+ The first of these is /o/ry. which is à measure for the curvature of the rotating surface.," The first of these is $l_0/r_\bnd$ , which is a measure for the curvature of the rotating surface."
+ In all, In all
+While a general solution to the initial value problem is not kuown. we use the particular solution obtained as follows: To solve the momentum constraints (33)) aud (31)) set. p.pefea=ea Oto leave padwpaw0 which may be satisfied by requizing pay=ec.,"While a general solution to the initial value problem is not known, we use the particular solution obtained as follows: To solve the momentum constraints \ref{ber-Hu}) ) and \ref{ber-hv}) ) set $p_x = p_z = \varphi,_a = \omega,_a = 0$ to leave $p_\Lambda \, \Lambda,_a -\, p_\Lambda,_a = 0$ which may be satisfied by requiring $p_\Lambda = c \,e^\Lambda$."
+" For sufficiently large ο, the Hamiltonian constraint nav be solved alecbraically for either p or r."," For sufficiently large $c$, the Hamiltonian constraint may be solved algebraically for either $p$ or $r$."
+ Iu general. this leaves as free data the four functions «c. 2. A. and either r or p.," In general, this leaves as free data the four functions $x$, $z$, $\Lambda$, and either $r$ or $p$."
+ Since there are four free functions at cach spatial point. we expect generic behavior.," Since there are four free functions at each spatial point, we expect generic behavior."
+ As a first case. we cousider polarized τς} models obtained by setting This condition is preserved umumerically as well as analytically.," As a first case, we consider polarized $U(1)$ models obtained by setting This condition is preserved numerically as well as analytically."
+ It has becu conjectured [11]. that polarized ((1) models are ΑΥΤΟ., It has been conjectured \cite{ber-bm2} that polarized $U(1)$ models are AVTD.
+ This is reasonable because the Mixiuaster poteutial-ike teri, This is reasonable because the Mixmaster potential-like term
+ This is reasonable because the Mixiuaster poteutial-ike terii, This is reasonable because the Mixmaster potential-like term
+In Fig.,In Fig.
+ 1. we present the LSD Stokes J and V obtained for ten orbital phases of 66 Eri., \ref{LSD} we present the LSD Stokes $I$ and $V$ obtained for ten orbital phases of 66 Eri.
+ The longitudinal magnetic field was determined from the first moment of theLSD Stokes V profile (?).., The longitudinal magnetic field was determined from the first moment of theLSD Stokes $V$ profile \citep{Kochukhov:2010}.
+ Due to the contribution of blends the continuum level of raw LSD Stokes 7 profiles is slightly offset from unity., Due to the contribution of blends the continuum level of raw LSD Stokes $I$ profiles is slightly offset from unity.
+ We use a constant factor to correct the continuum level., We use a constant factor to correct the continuum level.
+ This factor is applied to the LSD V and null profiles., This factor is applied to the LSD $V$ and null profiles.
+ Such a correction changes the value of the longitudinal magnetic field and its error by 2-3 G. For nine LSD V profiles we could measure the longitudinal field for each component separately., Such a correction changes the value of the longitudinal magnetic field and its error by 2–3 G. For nine LSD $V$ profiles we could measure the longitudinal field for each component separately.
+ At the phase 0.827 (see Fig. 1)), At the phase 0.827 (see Fig. \ref{LSD}) )
+ we have measured both components as a single star., we have measured both components as a single star.
+ The results of the longitudinal magnetic field measurements are presented in Table 1.., The results of the longitudinal magnetic field measurements are presented in Table \ref{tab1}.
+" Its columns contain the following information: the heliocentric Julian date of observation, the signal-to-noise ratio of the original spectra measured around 5200A,, the signal-to-noise ratio of the LSD V profiles, an orbital phase, determined from the binary solution discussed below."," Its columns contain the following information: the heliocentric Julian date of observation, the signal-to-noise ratio of the original spectra measured around 5200, the signal-to-noise ratio of the LSD $V$ profiles, an orbital phase, determined from the binary solution discussed below."
+" The next three columns represent measurements for the HgMn star (component A) as follows: the mean longitudinal magnetic field inferred from the LSD Stokes V, the longitudinal field inferred from the null profile, and a False Alarm Probability (FAP, see below)."," The next three columns represent measurements for the HgMn star (component A) as follows: the mean longitudinal magnetic field inferred from the LSD Stokes $V$, the longitudinal field inferred from the null profile, and a False Alarm Probability (FAP, see below)."
+" The last three columns give the same information for the component B. The errors of the magnetic field measurements were obtained by the standard error propagation, using the uncertainties of the LSD V profiles."," The last three columns give the same information for the component B. The errors of the magnetic field measurements were obtained by the standard error propagation, using the uncertainties of the LSD $V$ profiles."
+ The mean error of longitudinal field measurements is ~14 G for the HgMn component and ~20 G for the secondary., The mean error of longitudinal field measurements is $\sim$ 14 G for the HgMn component and $\sim$ 20 G for the secondary.
+ For the HgMn star the longitudinal field value reaches the confidence level of 2c only at the orbital phase 0.193., For the HgMn star the longitudinal field value reaches the confidence level of $\sigma$ only at the orbital phase 0.193.
+ For the secondary star the measurements of magnetic field do not exceed —1.5c level for most of the orbital phases., For the secondary star the measurements of magnetic field do not exceed $\sim$ $\sigma$ level for most of the orbital phases.
+" Thus, none of the individual mmeasurements indicate the presence of magnetic field in 66 Eri."," Thus, none of the individual measurements indicate the presence of magnetic field in 66 Eri."
+ Our magnetic field measurements for both components with their respective error bars are presented in Fig., Our magnetic field measurements for both components with their respective error bars are presented in Fig.
+ 2 as a function of orbital phase., \ref{Bz_phase} as a function of orbital phase.
+" The actual values of the longitudinal field measurements do not exceed the level of 30, so the periodic-like fluctuation of magnetic field measurements in both components are fully attributed to errors."," The actual values of the longitudinal field measurements do not exceed the level of $\sigma$, so the periodic-like fluctuation of magnetic field measurements in both components are fully attributed to errors."
+" Assuming a dipolar configuration of the magnetic field we made an assessment of the upper limit of its strength, which would be consistent with the entire set of mmeasurements."," Assuming a dipolar configuration of the magnetic field we made an assessment of the upper limit of its strength, which would be consistent with the entire set of measurements."
+" Using a relation between the longitudinal field and parameters of an oblique dipole (?),, we estimated the strength of dipolar magnetic field to be 60—70 G with an uncertainty of 50—70 G for both stars."," Using a relation between the longitudinal field and parameters of an oblique dipole \citep{Leroy:1994}, we estimated the strength of dipolar magnetic field to be 60–70 G with an uncertainty of 50–70 G for both stars."
+" However, based on the reduced x? statistics, the dipolar fit has no advantage over the null hypothesis."," However, based on the reduced $\chi^{2}$ statistics, the dipolar fit has no advantage over the null hypothesis."
+" At the same time, much more complex magnetic fields can exist in the atmospheres of both stars."," At the same time, much more complex magnetic fields can exist in the atmospheres of both stars."
+" To investigate possible presence of such complex magnetic structures, which might yield negligible(B,),, we employed the FAP analysis of LSD profiles."," To investigate possible presence of such complex magnetic structures, which might yield negligible, we employed the FAP analysis of LSD profiles."
+" FAP is a y? statistical estimate, which assesses if the Observed structure in the Stokes V line profile is produced by a random noise."," FAP is a $\chi^2$ statistical estimate, which assesses if the observed structure in the Stokes $V$ line profile is produced by a random noise."
+" Following the definition by ?,, these numbers"," Following the definition by \citet{Donati:1997}, , these numbers"
+reflection of the primary continuum from a photoionized gas or thermal emission from diffuse gas (MEKAL)). which was found to give good fits in a sample of Compton-thin Seyfert 2s observed by BeppoSAX (?)..,"reflection of the primary continuum from a photoionized gas or thermal emission from diffuse gas ), which was found to give good fits in a sample of Compton–thin Seyfert 2s observed by BeppoSAX \citep{risa02}."
+ The first one is parametrized (when self absorption effects are neglected) by a power law with the spectral index linked to that of the primary continuum responsible for the photoionization of the gas itself., The first one is parametrized (when self absorption effects are neglected) by a power law with the spectral index linked to that of the primary continuum responsible for the photoionization of the gas itself.
+ We further allowed for absorption by an intervening column density of neutral gas., We further allowed for absorption by an intervening column density of neutral gas.
+ The best fits for the two models are presented in Table 4:: even if marginally better from a statistical point of view. the fits have metal abundances unphysically low.," The best fits for the two models are presented in Table \ref{softfit}: even if marginally better from a statistical point of view, the fits have metal abundances unphysically low."
+ This means that the resulting spectrum lacks emission lines and is virtually indistinguishable from the already rejected bremsstrahlung interpretation., This means that the resulting spectrum lacks emission lines and is virtually indistinguishable from the already rejected bremsstrahlung interpretation.
+ Reflection from photoionized gas remains then the only tenable model., Reflection from photoionized gas remains then the only tenable model.
+" The observed intervening absorption is «2x107! em"". in excess of the Galactic one. and is possibly due to the host galaxy. which is seen almost edge-on."," The observed intervening absorption is $\simeq2\times 10^{21}$ $^{-2}$, in excess of the Galactic one, and is possibly due to the host galaxy, which is seen almost edge–on."
+ As expected from an origin from diffuse gas. the flux of this component is remarkably constant between the two XMM observations (see Table 4)). while the flux of the primary component almost doubles (see Sect. 2)).," As expected from an origin from diffuse gas, the flux of this component is remarkably constant between the two XMM observations (see Table \ref{softfit}) ), while the flux of the primary component almost doubles (see Sect. \ref{data}) )."
+ Therefore. the ratio," Therefore, the ratio"
+however. for 7)=0 the discrete spectrum disappears. together with the eigenfunctions.,"however, for $\eta=0$ the discrete spectrum disappears, together with the eigenfunctions."
+ Indeed. (his singularity in the induction operator lies al the heart of fast dynamo t(lieory (see.[orexample.Childress&Gilbert1995).," Indeed, this singularity in the induction operator lies at the heart of fast dynamo theory \citep[see, for example,][]{CG95}."
+. By contrast. the linear operator for the stability problem has a discrete spectrum of unstable modes and associated eisenfunctions. leading to an unambiguous determination of the electromotive lorce.," By contrast, the linear operator for the stability problem has a discrete spectrum of unstable modes and associated eigenfunctions, leading to an unambiguous determination of the electromotive force."
+" In this section we consider the emf resulting from the instability of a unidirectional mean magnetic field of the form: B=DBo(124-6z/d)e,. with non-zero values of the viscosity and magnetic diffusivitv."," In this section we consider the emf resulting from the instability of a unidirectional mean magnetic field of the form $\bfbarB=B_{0}(1+\zeta z/d)\mathbf{e}_{x}$, with non-zero values of the viscosity and magnetic diffusivity."
+" The instability evolves according to equations (16)) (22)) (with D,= 0). and the mean emf is calculated via horizontal averaging. given by (23)) (25)). and ihe weighting procedure (26))."," The instability evolves according to equations \ref{eq:u_pert}) \ref{eq:rho_pert}) ) (with $\Bbar_y =0$ ), and the mean emf is calculated via horizontal averaging, given by \ref{eq:Ex}) \ref{eq:Ez}) ), and the weighting procedure \ref{eq:weighted_emf}) )."
+ Figure 1. shows the three components of € plotted as functions of z. for a test case al colatitude 9=75°. representing the region in which most magnetic activity is observed on the solar surface.," Figure \ref{fig:emfs} shows the three components of $\bfcalE$ plotted as functions of $z$, for a test case at colatitude $\theta = 75^{\circ}$, representing the region in which most magnetic activity is observed on the solar surface."
+ For this parameter set. (he modes with largest growth rates have wave numbers (&./)=(0.70367.14.812) (with complex erowth rate s=0.082626— 0.0618231) for &>0. and (&./)=(—0.70334.14.312) (with s=0.082602+ 0.067830/) lor κ«0.," For this parameter set, the modes with largest growth rates have wave numbers $(k,l)=(0.70367,14.812)$ (with complex growth rate $s=0.082626-0.067823i$ ) for $k>0$, and $(k,l)=(-0.70334,14.812)$ (with $s=0.082602+0.067830i$ ) for $k<0$."
+ The ells in the figure are calculated rom the mode with largest growth rate: when there is a unique mode of maximum growth rate. as here. the differences between (he weighted aud unweighted emls are small.," The emfs in the figure are calculated from the mode with largest growth rate; when there is a unique mode of maximum growth rate, as here, the differences between the weighted and unweighted emfs are small."
+ It is of interest to note that the y-component of the emf pperpendicular to the imposed magnetic field) has the largest magnitude., It is of interest to note that the $y$ -component of the emf perpendicular to the imposed magnetic field) has the largest magnitude.
+ The calculations in indicate that in the absence of viscosity and magnetic diffusivitv. the ratio," The calculations in \\ref{sec:ideal} indicate that in the absence of viscosity and magnetic diffusivity, the ratio"
+distribution (Ratcliffeetal.1998).,distribution \citep{ratcliffe_IV_1998}.
+. Then we have probability density (9) of random variable & as follows 1n case of Gaussian distribution of velocities we have The convolution of c with e vields the distribution of pair numbers as a function of H., Then we have probability density $\Phi(\delta)$ of random variable $\delta$ as follows In case of Gaussian distribution of velocities we have The convolution of $\psi$ with $\Phi$ yields the distribution of pair numbers as a function of $\Pi$.
+ We used this to fit the observed. pair numbers in dilferent bins with the projected separation σ«Ly=2 Alpe alter randomized. background subtraction., We used this to fit the observed pair numbers in different bins with the projected separation $\sigma<L_{0}=2$ Mpc after randomized background subtraction.
+ The randomized. background. is caleulatecl as described in Sec. ??.., The randomized background is calculated as described in Sec. \ref{sec:2.7}.
+ Because the number of pairs is not large. there is a considerable variation of results for cilferent binnings. so we do not consider the numerical estimates as very reliable.," Because the number of pairs is not large, there is a considerable variation of results for different binnings, so we do not consider the numerical estimates as very reliable."
+" Llowever.we may argue that the model dependence on 5, and 5 parameters is rather insignificant."," However,we may argue that the model dependence on $\gamma_1$ and $\eta$ parameters is rather insignificant."
+" For illustration. we present the results. CLable 4) of fitting the pair numbers (alter randomized. background: subtraction) on account. of (11)) and (CIable 5) on account. of (12)) for two cdillerent binnings having minimal jackknife clispersion estiniates [or (63/7—(ut f23b7, "," For illustration, we present the results (Table 4) of fitting the pair numbers (after randomized background subtraction) on account of \ref{eq:zh3}) ) and (Table 5) on account of \ref{gauss}) ) for two different binnings having minimal jackknife dispersion estimates for $\langle v^2\rangle^{1/2}=\langle w^2/2\rangle^{1/2}$ ."
+Alnning | involves numbers. of pairs with line-of-sight. separations £7 from [five intervals 2.11]. 11.30]. 20.29]. 39.38]. 38.47] (Alpe).," Binning 1 involves numbers of pairs with line-of-sight separations $Pi$ from five intervals $[2,11]$ , $[11,20]$, $[20,29]$, $[29,38]$ , $[38,47]$ (Mpc)."
+ Binning 2 deals with separations from three intervals 2.17]. 17.32]. 32.41] (Alpe)," Binning 2 deals with separations from three intervals $[2,17]$, $[17,32]$, $[32,47]$ (Mpc)."
+ The bar charts in Fiο., The bar charts in Fig.
+ 16 and 17 we show corresponding numbers after the average backgrounc subtraction., \ref{fig_vel_eg1} and \ref{fig_vel_eg2} we show corresponding numbers after the average background subtraction.
+ Here we prefer to avoid separations less than 2 Alpe., Here we prefer to avoid separations less than 2 Mpc.
+ Also for illustration in Fig., Also for illustration in Fig.
+" 10 and 17 we presen re results of fitting for 5,=1.85 and dillerent. values of η.", \ref{fig_vel_eg1} and \ref{fig_vel_eg2} we present the results of fitting for $\gamma_1=1.85$ and different values of $\eta$.
+" As one can see [rom the results. the values of (ο215 do no 4vary considerably under rather cilferent. (even unrealistic) 'hoices of unknown parameters gj. 5,4: the variations are of 16 order of the jackknife dispersion estimate."," As one can see from the results, the values of $\langle v^2\rangle^{1/2}$ do not vary considerably under rather different (even unrealistic) choices of unknown parameters $\eta$, $\gamma_1$: the variations are of the order of the jackknife dispersion estimate."
+" In fact the values £02517 in Tables 5.66 do not represen real quasar velocity. dispersion. but a superposition of two components: ieο=ory,|5iA where τοσο ETEds the true velocity dispersion anc Όρων is due to the redshif measurement errors."," In fact the values $\langle v^2\rangle^{1/2}$ in Tables 6 do not represent real quasar velocity dispersion, but a superposition of two components: $\langle v^2\rangle = v_{err}^2 + \langle \bar{v}^2\rangle$, where $\langle\bar{v}^2\rangle^{1/2}$ is the true velocity dispersion and $v_{err}$ is due to the redshift measurement errors."
+ The last value corresponds to intrinsic emission lines shifts in quasars pointed out by Croometal. (2005).. as it is an estimation [rom emission-line or cross-correlation technique of redshift measurement caused by the impossibilitv of precise determination of the line centre due to its large width.," The last value corresponds to intrinsic emission lines shifts in quasars pointed out by \citet{croom_2005}, as it is an estimation from emission-line or cross-correlation technique of redshift measurement caused by the impossibility of precise determination of the line centre due to its large width."
+ For detailed investigation of this elfec see lDuchardsetal.(2002):Shenct(2007).," For detailed investigation of this effect see \citet{richards_2002,shen_2007}."
+. The redshifi measurement errors from the SDSS data base forour sample is A>=0.0019. which corresponds to 644ez/(1|2)230 kms. In this case we would obtain too large value of n2," The redshift measurement errors from the SDSS data base forour sample is $\Delta z=0.0019 $, which corresponds to $v _{err}\sim c\Delta z/(1+\bar{z})=230$ km/s. In this case we would obtain too large value of $\langle\bar{v}^2\rangle$."
+ Once work on the first version of the text had. been finished. the paper by Llewett&Wild(2010) was appeare with the error estimate ey.=600 km/s for SDSS data.," Once work on the first version of the text had been finished, the paper by \citet{hewett_10} was appeared with the error estimate $v_{err}=600$ km/s for SDSS data."
+ In this. case z-errors dominate. in. i£e7ο IMPΗΕ, In this case $z$ -errors dominate in ${\langle v^2\rangle}^{1/2}$.
+ we take the results from Table 5. which vield smaller jackknife dispersion estimate Lor (i2i1/2 than the residual. upper estimate of (73072«(6907GOOFA==340 km/s seems to he more adequate.," If we take the results from Table 5, which yield smaller jackknife dispersion estimate for $\langle{v}^2\rangle^{1/2}$, than the residual upper estimate of $\langle\bar{v}^2\rangle^{1/2}\le (690^2-600^2)^{1/2}=340$ km/s seems to be more adequate."
+ The values of the pairwise velocity dispersion. (7j=2073 corresponding to ‘Lables 5 are comparable with hat mentioned in previous studies of quasar clustering., The values of the pairwise velocity dispersion $\langle w^2\rangle= 2 \langle v^2\rangle$ corresponding to Tables 5 are comparable with that mentioned in previous studies of quasar clustering.
+ Outramοἱal.(2001) used the value 400 km/s obtained from he Llubble Volume Simulations: daAngelaetal.(2005) ixed the value SOO km/s and found that this value gives an adequate result for s«10 5.+ Alpe region ancl noted that it is dominated by the rms pairwise redshift error 600 kni/s: he same value SOO kis was used by daAnecla and the close value 600 km/s was used by Croom (2005)., \citet{outram_2001} used the value 400 km/s obtained from the Hubble Volume Simulations; \citet{daAngela_2005} fixed the value 800 km/s and found that this value gives an adequate result for $s<10$ $h^{-1}$ Mpc region and noted that it is dominated by the rms pairwise redshift error $600$ km/s; the same value 800 km/s was used by \citet{daAngela_2008} and the close value 690 km/s was used by \citet{croom_2005}.
+. Phe last value was chosen as a mean one of the range 630-750 kms: this is a quadrature superposition of (i) the real pairwise velocity dispersion recaleulatecl with redshift estimated [rom the galaxies. pairwise velocity dispersion 500 kms (Llawkinsetal. 2003).. (1) redshift’ measurement errorobtained. [rom repeat observations. ancl (iii) velocityerror due to intrinsic emission lines shifts in QSOs (Richards 2002)..," The last value was chosen as a mean one of the range 630-750 km/s; this is a quadrature superposition of (i) the real pairwise velocity dispersion recalculated with redshift estimated from the galaxies pairwise velocity dispersion 500 km/s \citep{hawkins_2003}, (ii) redshift measurement errorobtained from repeat observations, and (iii) velocityerror due to intrinsic emission lines shifts in QSOs \citep{richards_2002}. ."
+ Only Mountrichasetal.(2009). treating the velocity dispersion as a free parameter estimated.its value [or a sample of quasars and. luminous red. galaxies (LIC) as 620 km/s and 727 km/s (two values are the result. of, Only \citet{Mountrichas_2009} treating the velocity dispersion as a free parameter estimatedits value for a sample of quasars and luminous red galaxies (LRG) as 620 km/s and 727 km/s (two values are the result of
+Recurrent novae (RNe) are ordinary novae (binary svstems with mass accreting onto a white dwarl until (hermonuclear runaway is triggered) for which the recurrence time scale is between a decade and a century. such that more than one eruption has been observed 1964: Bode Evans 2008: Evans οἱ al.,"Recurrent novae (RNe) are ordinary novae (binary systems with mass accreting onto a white dwarf until thermonuclear runaway is triggered) for which the recurrence time scale is between a decade and a century, such that more than one eruption has been observed (Payne-Gaposchkin 1964; Bode Evans 2008; Evans et al."
+ 2008)., 2008).
+ To have the last recurrence time scale. the novae must have the white dwarl near the Chandrasekhar mass and have a high accretion rate.," To have the fast recurrence time scale, the novae must have the white dwarf near the Chandrasekhar mass and have a high accretion rate."
+ These properties. at lace value. imply that the white dwarf will soon exceed the Chandrasekhar mass and become a Type Ia supernova. and thus RNe are one of (he premier candidates for the progenitor class of these supernovae.," These properties, at face value, imply that the white dwarf will soon exceed the Chandrasekhar mass and become a Type Ia supernova, and thus RNe are one of the premier candidates for the progenitor class of these supernovae."
+ RNe t(vpically have relatively fast eruptions. high ejection velocities. aud small eruption ampliticdes when compared to ordinary novae.," RNe typically have relatively fast eruptions, high ejection velocities, and small eruption amplitudes when compared to ordinary novae."
+ Only ten RNe are known with certainty in our Milkv. Way (Schaefer 2010)., Only ten RNe are known with certainty in our Milky Way (Schaefer 2010).
+ U Scorpii (U Sco) previously erupted in March 1999 with a peak al V=7.5 mag (Schaeler 2010)., U Scorpii (U Sco) previously erupted in March 1999 with a peak at V=7.5 mag (Schaefer 2010).
+ In quiescence. it has Vzz17.6 and has deep/ofal eclipses taking it down to V—18.9 mag (Schaefer 2010) with an orbital period of 1.23 days (Schaeler 1990; Schaeler Ringwald 1995).," In quiescence, it has $\approx$ 17.6 and has deep eclipses taking it down to V=18.9 mag (Schaefer 2010) with an orbital period of 1.23 days (Schaefer 1990; Schaefer Ringwald 1995)."
+ U Sco is the lastest of all known novae. fading by three magnitudes from peak in just 2.6 davs. while its rise from minimum to peak is 6-12 hours (Schaefer 2010).," U Sco is the fastest of all known novae, fading by three magnitudes from peak in just 2.6 days, while its rise from minimum to peak is 6-12 hours (Schaefer 2010)."
+ No light echo was detected to deep limits alter the 1987 eruption (Schaefer 1983)., No light echo was detected to deep limits after the 1987 eruption (Schaefer 1988).
+ U Seo has now had ten known eruptions. in (he vears 1863. 1906. 1917. 1936. 1945. 1969. 1979. 1957. 1999 (Schaefer 2010). and now 2010 as we report in this paper.," U Sco has now had ten known eruptions, in the years 1863, 1906, 1917, 1936, 1945, 1969, 1979, 1987, 1999 (Schaefer 2010), and now 2010 as we report in this paper."
+ With the discovery of the 1917. 1945. and 1969 eruptions (Schaeler 2001: 2004). it has become apparent that U Sco has outbursts at intervals of 10+2 vears since 1900.," With the discovery of the 1917, 1945, and 1969 eruptions (Schaefer 2001; 2004), it has become apparent that U Sco has outbursts at intervals of $10\pm2$ years since 1900."
+ The exceptions to Chis are the two intervals of 19 and 24 vears. which are easily interpreted as being double intervals. with eruptions around 1927 and 1957 having been missed. (," The exceptions to this are the two intervals of 19 and 24 years, which are easily interpreted as being double intervals, with eruptions around 1927 and 1957 having been missed. ("
+U Sco is 3οἳ from the Sun every 28 November. so a significant [raction of its very [ast eruptions must be missed.),"U Sco is $3.6\degr$ from the Sun every 28 November, so a significant fraction of its very fast eruptions must be missed.)"
+ With Chis. it became apparent that the next eruption of U Seo should occur in (hie vear 2009--2.," With this, it became apparent that the next eruption of U Sco should occur in the year $2009\pm2$."
+ Schaefer (2005) made a better prediction. on the physical basis that (he time," Schaefer (2005) made a better prediction, on the physical basis that the time"
+as in the Milky Way.,as in the Milky Way.
+ The model is described in Matteucci et al. (, The model is described in Matteucci et al. (
+1998) where we address the reader for more details.,1998) where we address the reader for more details.
+" The star formation rate is given by: Le, normalized to the 1Utial total volume deusitv.", The star formation rate is given by: i.e. normalized to the initial total volume density.
+ c(t) 1s asstuued to drop to 0 at the ouset of the ealactic wind., $\psi(t)$ is assumed to drop to 0 at the onset of the galactic wind.
+ The quantity vis expressed in units of Cry| and represents the efficiency of star ornation. namely the iuverse of the time scale of star formation.," The quantity $\nu$ is expressed in units of $\rm Gyr^{-1}$ and represents the efficiency of star formation, namely the inverse of the time scale of star formation."
+ The star formation is ASSI to stop after t16 development of a galactic wine OCCULToο before than | Gar. from the begiuuiug of star formatic1. for all the ealaxies listed above.," The star formation is assumed to stop after the development of a galactic wind occurring before than 1 Gyr, from the beginning of star formation, for all the galaxies listed above."
+ Thereore. the star fornation rate iM these galaxies can be cojsdderec as a strong burst whic1i does not last more tha jid Cs a idis sjorter da more nassive svstenas.," Therefore, the star formation rate in these galaxies can be considered as a strong burst which does not last more than 1 Gyr and is shorter in more massive systems."
+" This is obtaiucc by assuwine that the star formation efficiency dICLCASCS with eaactic mass tlis producing au "" inverse wind OFfect. as ¢escribed in Matteucci (1991). where the ealactic wind occtis before i more massive than iu less massive ¢5lipticaS."," This is obtained by assuming that the star formation efficiency increases with galactic mass thus producing an “ inverse wind” effect, as described in Matteucci (1994), where the galactic wind occurs before in more massive than in less massive ellipticals."
+ As a consequence of this. the star formation )eriod is onuecr m smaller svsteis thus allowing the SNe Ta to subsautially pollute the ISAL," As a consequence of this, the star formation period is longer in smaller systems thus allowing the SNe Ia to substantially pollute the ISM."
+ This effect can explain 1ο observed increase of the [a/Fe] ratio with galactic nass (AVGxthev et al., This effect can explain the observed increase of the $\alpha$ /Fe] ratio with galactic mass (Worthey et al.
+ 1992: Matteucci 1991). which is rot obtaired in oa classic wind scenario. where the more nassive objects formi stars for a longer time (Larson. 1971).," 1992; Matteucci 1994), which is not obtained in a classic wind scenario, where the more massive objects form stars for a longer time (Larson, 1974)."
+ The galactic wind develops as à consequence of the enerev transfer froii SNe into the ISM., The galactic wind develops as a consequence of the energy transfer from SNe into the ISM.
+" In fact. when the hermal enerevOo, of the Ooeas becomes IlarecrOo than the bindingC» cnerev of the eas. the wind starts (Arimoto Yoshii. 1987: Matteucci Tornambe 1987: Matteucci 1992. 199 Dipi loc al."," In fact, when the thermal energy of the gas becomes larger than the binding energy of the gas, the wind starts (Arimoto Yoshii, 1987; Matteucci Tornambe' 1987; Matteucci 1992, 1994; Pipino et al."
+ 2002)., 2002).
+ In order to compute tlic|] udine enerev of the e:Us solic asstuuptions have to be uade about he ealactic potential wel., In order to compute the binding energy of the gas some assumptions have to be made about the galactic potential well.
+ In particular. it is ussted that a1 cliplcas possess heavy hit diffuse dark 11aatter halos: : ratio heween the haf-lish radius aud t1ο radius of li dark luatter core {ηνμονRdar το. aud a ratio dar tohWUnOUs mass of ]0 are assunied.," In particular, it is assumed that all ellipticals possess heavy but diffuse dark matter halos; a ratio between the half-light radius and the radius of the dark matter core $R_{luminous}/R_{dark}$ =0.10 and a ratio dark to luminous mass of 10 are assumed."
+ The D aud I& Iuniuosities for the ellipticals of differeut lle9SCR are colmputed by means of the spectro-photomoetric inodel o| Jimenez et al. (, The B and K luminosities for the ellipticals of different masses are computed by means of the spectro-photometric model of Jimenez et al. (
+1999) and are usxl to compute the SN Ta vate iun SM aud the nova rate per unit of Ly. respectively.,"1999) and are used to compute the SN Ia rate in SNu and the nova rate per unit of $L_K$, respectively."
+ The mai cliffereices between the inodel for the Galaxy. and the modcl for an elliptical οalaxy couceru the different SER. which is much strougcr in the carliest stages. and then is set o zero after the ealacic Wind in the case of ellipticals.," The main differences between the model for the Galaxy, and the model for an elliptical galaxy concern the different SFR, which is much stronger in the earliest stages, and then is set to zero after the galactic wind in the case of ellipticals."
+" It is worth notius that «ealactic winds in the case of elliptica enlaxies seen necessary to explain their lack of gas aud tιο Chemica παςit of the ICM. wherCas he Galactic model does not take iuto account the occurrence of a stroug πια, maiilv. because of the strong eyavitatknal potential well associated witji our Galaxy. aud the xeseuce of the gas in the Galactic clik."," It is worth noting that galactic winds in the case of elliptical galaxies seem necessary to explain their lack of gas and the chemical enrichment of the ICM, whereas the Galactic model does not take into account the occurrence of a strong wind, mainly because of the strong gravitational potential well associated with our Galaxy, and the presence of the gas in the Galactic disk."
+ We adopt two differeut initia mass functions (IAIFs): the Salpeter Grpapr1.35 ) one aud the Arimoto Yoshii (OST) ΕΕ 0.95) One., We adopt two different initial mass functions (IMFs): the Salpeter $x_{IMF}=1.35$ ) one and the Arimoto Yoshii (1987) $x_{IMF}=0.95$ ) one.
+ Iu tevet. successful models of chemical evolution of ellipticals have shown," In fact, successful models of chemical evolution of ellipticals have shown"
+Figue 2 shows the Eddington cnhaucement factor for variable inclination. of slabs (upper panels) and for changing density contrast of vertical slabs (lower panels).,Figure 2 shows the Eddington enhancement factor for variable inclination of slabs (upper panels) and for changing density contrast of vertical slabs (lower panels).
+ Iu all cases. the Thomsou depth of the Ligh density slabs Tj ls constant but the optical depth across low density regions mereases from values7z« οτι >1 from left to right (see caption of Fie.," In all cases, the Thomson depth of the high density slabs $\tau_{h}$ is constant but the optical depth across low density regions increases from values$\tau_{l}<1$ to $\tau_{l}>1$ from left to right (see caption of Fig."
+ 2 for more details)., 2 for more details).
+ As expected. the Eddington factor increases as the slabs rotate toward. the vertical direction because the atinosphere effectively becomes more porous (see upper panels}.," As expected, the Eddington factor increases as the slabs rotate toward the vertical direction because the atmosphere effectively becomes more porous (see upper panels)."
+ When the slabs are vertical. the flux cnhancemenut factor increases as the density contrast σι becomes larger for constant mean density.," When the slabs are vertical, the flux enhancement factor increases as the density contrast $\sigma_{h}/\sigma_{l}$ becomes larger for constant mean density."
+ This is due to the fact that the vohune filling factor of the high-density eas decreases while that of the low-density eas increases. but the respective masses of the two density phases remain the same.," This is due to the fact that the volume filling factor of the high-density gas decreases while that of the low-density gas increases, but the respective masses of the two density phases remain the same."
+" Therefore. the mean. voluueaveielted. flux is where fe is the volume filius factor of the dense gas and Fy aud F), are the fluxes propagating through tenuous and dense regions. respectively."," Therefore, the mean, volume-weighted, flux is where $f_{v}$ is the volume filling factor of the dense gas and $F_{l}$ and $F_{h}$ are the fluxes propagating through tenuous and dense regions, respectively."
+" As the density coutrast lnereases and ff, decreases. radiation tends to flow primarily through the low deusity channels aud. therefore. tore flux is necessary to exert the same total force as in the homogeneous case because radiation iuteracts less efficicutly with teuuous eas."," As the density contrast increases and $f_{v}$ decreases, radiation tends to “flow” primarily through the low density channels and, therefore, more flux is necessary to exert the same total force as in the homogeneous case because radiation interacts less efficiently with tenuous gas."
+ Quautitatively. iu the diffusion nuit. we have (Shaviv1998) When most volume is in the low-density phase but most mass is in the ligh-deusity phase (ie. fe>0 or €n/Q> O) then |zfopy/pi it fe9ppp.," Quantitatively, in the diffusion limit, we have \citep{sh98}
+ When most volume is in the low-density phase but most mass is in the high-density phase (i.e., $f_{v}\rightarrow 0$ or $\xi_{h}/\xi_{l}\rightarrow 0$ ), then $l\approx f_{v}\rho_{h}/\rho_{l}$ if $f_{v}\gg\rho_{l}/\rho_{h}$."
+ This qualitatively explains why 7 decreases with 7= lat constant density contrast 05/0; (cf, This qualitatively explains why $l$ decreases with $\tau_{l}\ge 1$ at constant density contrast $\sigma_{h}/\sigma_{l}$ (cf.
+ third aud fourth columns on Fie., third and fourth columns on Fig.
+ 2)., 2).
+ At small optical depth τι. equation (23) would lead to very inaccurate answers.," At small optical depth $\tau_{l}$, equation (23) would lead to very inaccurate answers."
+" For example. equation (23) predicts |~23 for vertical slabs with 75=(kl aud e,/o;=100. compared to the actual value 7—5 aud our analytic result —L from eq. ("," For example, equation (23) predicts $l\sim 23$ for vertical slabs with $\tau_{l}=0.1$ and $\sigma_{h}/\sigma_{l}=100$, compared to the actual value $l\sim 5$ and our analytic result $\sim 4$ from eq. ("
+21) (cf,21) (cf.
+ lower left panel), lower left panel).
+" This discrepancy is due larecly to ueelect of the amisotropy of the radiation feld. whereas our approach eives much more accurate results even in such au extreme case,"," This discrepancy is due largely to neglect of the anisotropy of the radiation field, whereas our approach gives much more accurate results even in such an extreme case."
+ Moreover. note that the “anisotropy term” in our expression for the flux cuhancement factor. which ix proportional to (No)?=(ojσι)”. vanishes for large Thomson depths aud thus equatious (21) aud (15) reduce to equation (23) iu the diffusion limit.," Moreover, note that the “anisotropy term” in our expression for the flux enhancement factor, which is proportional to $(\Delta\sigma)^{2}
+=(\sigma_{h}-\sigma_{l})^{2}$, vanishes for large Thomson depths and thus equations (21) and (15) reduce to equation (23) in the diffusion limit."
+" We also considered ""unulti-streaua approximation schemes in order to account for the radiatiou anisotropy. but found the ""iutensitv moment” approach developed here to be in significantly better agreement with Monte Carlo simulations."," We also considered “multi-stream” approximation schemes in order to account for the radiation anisotropy, but found the “intensity moment” approach developed here to be in significantly better agreement with Monte Carlo simulations."
+ We have considered radiative transfer deep within extremely inhomogeneous atinospheres. and lave demonstrated that. under such conditions. radiation viscosifv — Le. the off-diagonal clemeuts of the radiation stress tensor — plays an maportaut role.," We have considered radiative transfer deep within extremely inhomogeneous atmospheres, and have demonstrated that, under such conditions, radiation viscosity – i.e., the off-diagonal elements of the radiation stress tensor – plays an important role."
+ Our approach is sienificautly more accurate than approaches based on the diffusion οςation aud iulti-trezu approximation., Our approach is significantly more accurate than approaches based on the diffusion equation and multi-stream approximation.
+ The technique develyped here cau be applied to the noulinear evohtion of radiation-driven dustabilifies iun accretion clisks., The technique developed here can be applied to the nonlinear evolution of radiation-driven instabilities in accretion disks.
+" Iu parictlar. it can be used to study the dynamical coupine of iater and radiation iu order to determune the characteristic leneth scales and density contrasts arising frou ""photoilu ibble instability."," In particular, it can be used to study the dynamical coupling of matter and radiation in order to determine the characteristic length scales and density contrasts arising from “photon bubble” instability."
+ This. inturn. will permit a sel—COUSISenu determination of the magnitude of the Eddiietol csnhancement factor im radiatiou-donünated accretion disks.," This, in turn, will permit a self-consistent determination of the magnitude of the Eddington enhancement factor in radiation-dominated accretion disks."
+ We also sugeest that our method could be ΠιοΟΥ)oratec nno radiation hvdrodynauies codes such as the RIID modue for ZEUS (TurneraudStone2001).., We also suggest that our method could be incorporated into radiation hydrodynamics codes such as the RHD module for ZEUS \citep{tu01}. .
+ This work was supported in part bv NSF eraut. AST, This work was supported in part by NSF grant AST--9876887.
+ Wef hank Nix Shaviv aud Neal Turuer for conmunenuts on the nauuscript., We thank Nir Shaviv and Neal Turner for comments on the manuscript.
+altitude 5. and the distance from the observer r. separately or the Northern and Southern Galactic hemispheres.,"latitude $b$, and the distance from the observer $r$, separately for the Northern and Southern Galactic hemispheres."
+ We determined: selection. functions For clusters populating rich and poor superclusters. using a threshold richness of Noy S.," We determined selection functions for clusters populating rich and poor superclusters, using a threshold richness of $N_{cl}=8$ ."
+ Phe influence of the choice of the threshold richness Αι gawll be discussed in the next Section., The influence of the choice of the threshold richness $N_{cl}$ shall be discussed in the next Section.
+ In Figure 1. we show the results of the determination of 1ο selection function for clusters of galaxies with measured redshifts., In Figure 1 we show the results of the determination of the selection function for clusters of galaxies with measured redshifts.
+ Phe number of clusters vs. the Galactic latitude was determined. as a function of sind., The number of clusters vs. the Galactic latitude was determined as a function of $\sin b$.
+ Differences between 10 (wo hemispheres are small. thus in Figure 1 we present 1e mean of both hemispheres.," Differences between the two hemispheres are small, thus in Figure 1 we present the mean of both hemispheres."
+ Data are normalised to unit density at sind=1., Data are normalised to unit density at $\sin b =1$.
+ We see an almost linear decrease of the number density of clusters with sinb., We see an almost linear decrease of the number density of clusters with $\sin b$.
+ This linear regression. D(b)=(snb—sinóu)/(1snb). is given by the value Sy=sinbdy where the density of cluster reaches 0. and. it was used to calculate Poisson samples for the correlation function.," This linear regression, $D(b)=(\sin b - \sin b_0)/(1-\sin b_0)$, is given by the value $s_0=\sin b_0$ where the density of cluster reaches 0, and it was used to calculate Poisson samples for the correlation function."
+ To determine the distance dependence of the selection function the spatial density. of clusters of galaxies was calculated in) concentric spherical shells of thickness 20Alpe. for each hemisphere separately.," To determine the distance dependence of the selection function the spatial density of clusters of galaxies was calculated in concentric spherical shells of thickness 20, for each hemisphere separately."
+ Fluctuations are rather large. thus for this sample of clusters the mean regression was derived for both hemispheres.," Fluctuations are rather large, thus for this sample of clusters the mean regression was derived for both hemispheres."
+ Phe spatial density can be represented by a linear law: DOr)=dy αι). where dy and dj are constants. ancl (ry ds the outer radius of the sample.," The spatial density can be represented by a linear law: $D(r)=d_0-d_1(r/r_1)$ , where $d_0$ and $d_1$ are constants, and $r_1$ is the outer radius of the sample."
+ Values of the selection function xwameters dy. and. di. found. for various. subsamples of clusters. are given in Table 1.," Values of the selection function parameters $d_0$, and $d_1$, found for various subsamples of clusters, are given in Table 1."
+ A similar analvsis of the selection. function. was mace or the sample of all L304 clusters., A similar analysis of the selection function was made for the sample of all 1304 clusters.
+" Here. too. the sample was divided into high- and low-density populations using he same threshold NV,=8."," Here, too, the sample was divided into high- and low-density populations using the same threshold $N_{cl}=8$."
+ Table 1 shows that parameters of the distance dependence in the Northern and Southern yemisphere (denoted with subscripts N and S. respectively) are identical in most. cases.," Table 1 shows that parameters of the distance dependence in the Northern and Southern hemisphere (denoted with subscripts $N$ and $S$, respectively) are identical in most cases."
+ Only the cluster sample of all clusters in low-density regions is large enough to determine xwameters of the distance dependence separately for both Yemispheres., Only the cluster sample of all clusters in low-density regions is large enough to determine parameters of the distance dependence separately for both hemispheres.
+ Here dov is smaller than dos. which reflects the act that the number-density of the Northern cluster sample is lower than that of the Southern one.," Here $d_{0N}$ is smaller than $d_{0S}$, which reflects the fact that the number-density of the Northern cluster sample is lower than that of the Southern one."
+ Parameters for the selection elfect in Galactic latitucle are similar for the sample of all clusters and that of clusters with measured. redshifts., Parameters for the selection effect in Galactic latitude are similar for the sample of all clusters and that of clusters with measured redshifts.
+ In this Section we discuss the correlation. function. of AbellACO clusters of galaxies in. various environments., In this Section we discuss the correlation function of Abell–ACO clusters of galaxies in various environments.
+ As noted. above. clusters in) high-density environment (rich superclusters) form a fairly regular three-dimoensional network. whereas clusters in low-density environment (isolated: clusters and. clusters in. poor and medium rich," As noted above, clusters in high-density environment (rich superclusters) form a fairly regular three-dimensional network, whereas clusters in low-density environment (isolated clusters and clusters in poor and medium rich"
+positions are in good agreement with those derived from the slit spectra.,positions are in good agreement with those derived from the long-slit spectra.
+" For knots A and N, the ratios obtained from the IFS data are in better agreement with those derived by Rieraetαἱ(20032) than with those derived by Reipurth,Raga&Heathcote(1996);; (ii) IFS line-ratio maps revealed some regions in which the degree of gas excitation is low."," For knots A and N, the ratios obtained from the IFS data are in better agreement with those derived by \citet{Rie03a} than with those derived by \citet{Rei96}; (ii) IFS line-ratio maps revealed some regions in which the degree of gas excitation is low."
+" In particular, the highest 11]//Ha ratios (> 1.5) are found around knot H and slightly lower ratios, which are also compatible with a low degree of excitation, are found around knots R and S. This wider range of excitation degree could not be detected by previous long-slit observations, which did not sample these low excitation regions; (iii) no clear systematic trend in the spatial distribution of the excitation was found neither along the jet outflow nor across the jet beam."," In particular, the highest $\alpha$ ratios $\ge$ 1.5) are found around knot H and slightly lower ratios, which are also compatible with a low degree of excitation, are found around knots R and S. This wider range of excitation degree could not be detected by previous long-slit observations, which did not sample these low excitation regions; (iii) no clear systematic trend in the spatial distribution of the excitation was found neither along the jet outflow nor across the jet beam."
+" However, the integrated Π]//Ηα and rr]//Ha: maps indicate that the excitation is higher for the northern jet knots A to D (where the outflow has a higher collimation) than beyond knot E, from the region where the outflow widens and the knots begin to lose their alignment relative to the axis direction traced by knots A-C. For this jet region ffrom knots E to Q) the line ratio maps show some signatures of an increment in the excitation moving outwards, from the knot peak to the eastern knot border and to the interknot."," However, the integrated $\alpha$ and $\alpha$ maps indicate that the excitation is higher for the northern jet knots A to D (where the outflow has a higher collimation) than beyond knot E, from the region where the outflow widens and the knots begin to lose their alignment relative to the axis direction traced by knots A-C. For this jet region from knots E to Q) the line ratio maps show some signatures of an increment in the excitation moving outwards, from the knot peak to the eastern knot border and to the interknot."
+ The line ratio maps are thus suggestive of a decrease in the excitation from east to west across the jet beam., The line ratio maps are thus suggestive of a decrease in the excitation from east to west across the jet beam.
+" This is true for all knots, except around knot I (at the western jet border), where the excitation is higher and similar to that derived in the northern (A to D) knots."," This is true for all knots, except around knot I (at the western jet border), where the excitation is higher and similar to that derived in the northern (A to D) knots."
+ The spatial distribution of the electron density (ιο) is traced by the 66716/6731 map., The spatial distribution of the electron density $n_\rmn{e}$ ) is traced by the 6716/6731 map.
+ Figure 3 shows the map created from the IFU data (with the assumptions mentioned before for the excitation maps)., Figure \ref{cfi} shows the map created from the IFU data (with the assumptions mentioned before for the excitation maps).
+" At the weakest knots, the estimated uncertainty is15%."," At the weakest knots, the estimated uncertainty is."
+". The values found range from ~ 0.9 to ~ 1.4, which correspond to ne ranging from 1000 to 50 cm? respectively (these values were derived using the task of the package, and assuming Το=10 K)."," The values found range from $\sim$ 0.9 to $\sim$ 1.4, which correspond to $n_\rmn{e}$ ranging from 1000 to 50 $^{-3}$ respectively (these values were derived using the task of the package, and assuming $T_\rmn{e}=10^4$ K)."
+" The electron density cannot be properly evaluated around the knot R region, because the derived line ratios are higher than the values for which the 66716/6731 is density sensitive in the low-density limit."," The electron density cannot be properly evaluated around the knot R region, because the derived line ratios are higher than the values for which the 6716/6731 is density sensitive in the low-density limit."
+" In general, the trend found is a decrease in the ne, moving along the jet from north to south."," In general, the trend found is a decrease in the $n_\rmn{e}$, moving along the jet from north to south."
+" This result is consistent with the density behaviour of the knots derived from the long-slit spectra by Reipurth,Raga&Heathcote(1996) and Rieraetal(2003a).", This result is consistent with the density behaviour of the knots derived from the long-slit spectra by \citet{Rei96} and \citet{Rie03a}.
+". However, the spatial distribution of the density obtained through IFS refcfi) is more complex than that outlined from long-slit spectroscopy, and could not be properly described based on this kind of observations."," However, the spatial distribution of the density obtained through IFS \\ref{cfi}) ) is more complex than that outlined from long-slit spectroscopy, and could not be properly described based on this kind of observations."
+" The density map indicates that, in general, the electron densities are higher around the knot peak positions and decrease towards the edges of the knots and towards the interknot gas."," The density map indicates that, in general, the electron densities are higher around the knot peak positions and decrease towards the edges of the knots and towards the interknot gas."
+" Across the jet beam, a trend is observed in the knots peaking towards the west side of the axis defined by knots A-C kknots E, I, N, P) that are denser than knots peaking towards the east side of this axis kknots D, J, O, Q)."," Across the jet beam, a trend is observed in the knots peaking towards the west side of the axis defined by knots A–C knots E, I, N, P) that are denser than knots peaking towards the east side of this axis knots D, J, O, Q)."
+ The density map also shows some local departures of this general trend that are significant., The density map also shows some local departures of this general trend that are significant.
+" In particular, it should be noted that the highest n. is reached around knot E ~ 40 arcsec south from knot A)."," In particular, it should be noted that the highest $n_\rmn{e}$ is reached around knot E $\sim$ 40 arcsec south from knot A)."
+" In fact, the IFU data revealed an enhancement in the electron density, reaching ne values of ~ 1000 επιὃ and with a maximum of ~ 1300 cm? located ~ 3 arcsec southwest of the knot E peak intensity."," In fact, the IFU data revealed an enhancement in the electron density, reaching $n_\rmn{e}$ values of $\simeq$ 1000 $^{-3}$ and with a maximum of $\simeq$ 1300 $^{-3}$ located $\simeq$ 3 arcsec southwest of the knot E peak intensity."
+" This denser clump, extending over ~ 30 arcsec”, appears well delimited at its eastern edge, as shown by sharp fall of ne, which reaches values of ~ 200-300 cm? by aonly moving ~ 2 arcsec east from the E peak intensity."," This denser clump, extending over $\simeq$ 30 $^2$, appears well delimited at its eastern edge, as shown by a sharp fall of $n_\rmn{e}$, which reaches values of $\simeq$ 200–300 $^{-3}$ by only moving $\simeq$ 2 arcsec east from the E peak intensity."
+" In addition, although n. gently decreases from knots A (~ 600 cm?) to H (~ 400 cm""), the knot I, southwest of all of these knots, has a ne similar to knot B (~ 500 cm?), while knot D, southeast of knot C, has similar low values (~ 150 cm?) to those found at the southern, more rarefied knots, beyond knot N. The emission lines at most of the positions mapped in HH 110 show asymmetric profiles that cannot be properly fitted by a single Gaussian."," In addition, although $n_\rmn{e}$ gently decreases from knots A $\sim$ 600 $^{-3}$ ) to H $\sim$ 400 $^{-3}$ ), the knot I, southwest of all of these knots, has a $n_\rmn{e}$ similar to knot B $\sim$ 500 $^{-3}$ ), while knot D, southeast of knot C, has similar low values $\sim$ 150 $^{-3}$ ) to those found at the southern, more rarefied knots, beyond knot N. The emission lines at most of the positions mapped in HH 110 show asymmetric profiles that cannot be properly fitted by a single Gaussian."
+" Hence, we calculated the flux-weighted mean radial velocity (or first-order intensity momentum) instead of the line centroids of a Gaussian fit, and the flux-weighted rms width of the line (or second-order intensity momentum), proportional to the FWHM of the line, and described the integrated velocity field of the from these intensity momenta."," Hence, we calculated the flux-weighted mean radial velocity (or first-order intensity momentum) instead of the line centroids of a Gaussian fit, and the flux-weighted rms width of the line (or second-order intensity momentum), proportional to the FWHM of the line, and described the integrated velocity field of the from these intensity momenta."
+ HH 110 shows a velocity field behaviour very similar in all the mapped lines refvel))., HH 110 shows a velocity field behaviour very similar in all the mapped lines \\ref{vel}) ).
+" The radial velocities are blueshifted (Visr ranging from c -10 to -80 iin Ho), with ankm increase towards more blueshifted values moving from north to south."," The radial velocities are blueshifted $\mbox{\vlsr}$ ranging from $\sim$ –10 to –80 in $\alpha$ ), with an increase towards more blueshifted values moving from north to south."
+ Velocities that were derived from Ha are blueshifted by ~ 20 relative to those derived from the and ((the velocities derived from these two lines are similar)., Velocities that were derived from $\alpha$ are blueshifted by $\sim$ 20 relative to those derived from the and (the velocities derived from these two lines are similar).
+" The velocity offset found between Ho and lis not an artifact caused by an inaccurate wavelength calibration of the IFS data, since it has been also found in long-slit, higher resolution data citealpLop05))."," The velocity offset found between $\alpha$ and is not an artifact caused by an inaccurate wavelength calibration of the IFS data, since it has been also found in long-slit, higher resolution data \\citealp{Lop05}) )."
+" A more detailed inspection of the velocity maps reveals a more complex structure of the velocity field: (1) the less blueshifted velocities (Vis » 10 to 15 iin Ha) are found around the northern emissions X, Y and in knot"," A more detailed inspection of the velocity maps reveals a more complex structure of the velocity field: (i) the less blueshifted velocities $\mbox{\vlsr}$ $\sim$ –10 to –15 in $\alpha$ ) are found around the northern emissions X, Y and in knot"
+fields. discussed here. are oOgiven in asterisks.,"fields, discussed here, are given in asterisks."
+ In the rightoO plot. ap against frequency is displaved.," In the right plot, $\alpha_P$ against frequency is displayed."
+ The connected points are from the same area observed at different frequencies| frou 108. MITIZz to 1411. MIIz (in the Droww Spoclstra paper)., The connected points are from the same area observed at different frequencies from 408 MHz to 1411 MHz (in the Brouw Spoelstra paper).
+ The first couclusion from Fig., The first conclusion from Fig.
+ 12. is that multipole spectra indices vary over the slo froma z1 to 3. without showing a clear correlation with Galactic longitude. while oulv the WENS: subfields show a dependence of spectral ludex 1 latitude.," \ref{f9:lit} is that multipole spectral indices vary over the sky from $\alpha \approx 1$ to 3, without showing a clear correlation with Galactic longitude, while only the WENSS subfields show a dependence of spectral index on latitude."
+ However. all surveys lave been done at differcut resolutious aud frequencies. aud the regious used to compute power spectra are of differeut sizes.," However, all surveys have been done at different resolutions and frequencies, and the regions used to compute power spectra are of different sizes."
+ This cau explain why a possible dependence of a on latitude was not clearly seen in the other studies., This can explain why a possible dependence of $\alpha$ on latitude was not clearly seen in the other studies.
+ The large variation in slopes of augular power spectra in P? iudicates that interpretation of the slope is not straightforward. possibly due to large influence of depolarization mechanisms.," The large variation in slopes of angular power spectra in $P$ indicates that interpretation of the slope is not straightforward, possibly due to large influence of depolarization mechanisms."
+ Care uust therefore be taken in extrapolating the results to ueher frequencies., Care must therefore be taken in extrapolating the results to higher frequencies.
+ Furthermore. spectral mdices show an merease with requeucy from 105 MIIZ o 1.1 GIIz.," Furthermore, spectral indices show an increase with frequency from 408 MHz to 1.4 GHz."
+ This ποσα» that the oower spectra become steeper. so the relative amount of simudlscale structure decreases.," This means that the power spectra become steeper, so the relative amount of small-scale structure decreases."
+ This couk be due to the arec Faraday rotation at low frequencies., This could be due to the large Faraday rotation at low frequencies.
+ Typical Ax of: 5 aare present in the Brown Spoclstra data (Spoclstra 1981). and will rotate polarization angles at 325 MIIz by about250°.," Typical $RM$ s of 5 are present in the Brouw Spoelstra data (Spoelstra 1984), and will rotate polarization angles at 325 MHz by about."
+. Variations in RAL of a few eeive angle variations of over907.. which would cause bean depolarization if the angle variatious occur on scales sinaller than the beam Hu this case)," Variations in $RM$ of a few give angle variations of over, which would cause beam depolarization if the angle variations occur on scales smaller than the beam in this case)."
+ Beam depolarization only acts on scales of the svuthesized beam. aud therefore creates additional structure ou small scales in P. which flattens the power spectrum.," Beam depolarization only acts on scales of the synthesized beam, and therefore creates additional structure on small scales in $P$, which flattens the power spectrum."
+ A ARAL of 5 wwould cause a variation iu polarization angle of about aat 820 ΑΠΣ. and of no more than aat 1.1 GIIz.," A $\Delta RM$ of 5 would cause a variation in polarization angle of about at 820 MHz, and of no more than at 1.4 GHz."
+ So at frequencies above 1.41 GIIz. à ARA of Ὁ wwould cause negheible beam depolarization.," So at frequencies above 1.4 GHz, a $\Delta RM$ of 5 would cause negligible beam depolarization."
+ In addition. the resolution of the observations ecnerally imereases with increasing frequency. which would also cause a decrease i beam depolarization.," In addition, the resolution of the observations generally increases with increasing frequency, which would also cause a decrease in beam depolarization."
+ This might explain why above ee.l GHz the spectral index does uot appear to be correlated with frequency., This might explain why above 1.4 GHz the spectral index does not appear to be correlated with frequency.
+ The fact that spectral iudices in the two WSRT reeions are much higher than would be sxpected from this arguiueut can e due to the criteria sed to select the two fields., The fact that spectral indices in the two WSRT regions are much higher than would be expected from this argument can be due to the criteria used to select the two fields.
+ The disadvantage of using augular power spectra is that a regular erid of data is required., The disadvantage of using angular power spectra is that a regular grid of data is required.
+ If the data are very nreeululv spaced in the case of data frou pulsars or extragalactic poiut sources). if is better to use the structure fiction which in principle gives the same information. but can be caleulated casily for nregululv spaced data.," If the data are very irregularly spaced in the case of data from pulsars or extragalactic point sources), it is better to use the structure function which in principle gives the same information, but can be calculated easily for irregularly spaced data."
+ The structure function SF of a radiation field IX. as a fiction of distance lag d. is where VOr;) is the value of ficld (X. at position e; aud is the number of data points.," The structure function $SF$ of a radiation field $X$, as a function of distance lag $d$, is where $X(x_i)$ is the value of field $X$ at position $x_i$ and $N$ is the number of data points."
+ If the power spectra. of RAL is a power law with spectral index a. then the structure function SFpa ds (Simonetti ct 14398D).," If the power spectrum of $RM$ is a power law with spectral index $\alpha$, then the structure function $SF_{RM}$ is (Simonetti et 1984)."
+ We determined the structure functions of RAL to compare wihn existing cmates of the structure function of Calacic RAM frou polarize extragalactic point sources., We determined the structure functions of $RM$ to compare with existing estimates of the structure function of Galactic $RM$ from polarized extragalactic point sources.
+ As the deterwination of structure functions does not recmire a reguar erid. we can compute the SF including aud excluding “bad” data points to examine how the structure functions of P iux RAL change.," As the determination of structure functions does not require a regular grid, we can compute the $SF$ including and excluding “bad” data points to examine how the structure functions of $P$ and $RM$ change."
+ This will allow us to estimate the effect of “bad” data in the power spectra in P and RAL compute iu Sect. 3., This will allow us to estimate the effect of “bad” data in the power spectra in $P$ and $RM$ computed in Sect. \ref{s9:ps}.
+ Structure functions SEga; in the Aurigaand Worologimm regions are plotted against distance d iu degrees in the log-log plots iu Fig., Structure functions $SF_{\scriptsize RM}$ in the Aurigaand Horologium regions are plotted against distance $d$ in degrees in the log-log plots in Fig.
+ 13) (solid line). where error bars denote the standard deviation.," \ref{f9:sf_rm} (solid line), where error bars denote the standard deviation."
+ The minimum distauce shown is dxτι," The minimum distance shown is $d
+\approx$."
+", For the evaluation of the SE. ouly ~rehably cternuned” RAL values are used. according to he definition in Sect. 2.1.."," For the evaluation of the SF, only “reliably determined” $RM$ values are used, according to the definition in Sect. \ref{s9:wsrtdata}."
+ Although the spectruni is consistent with a flat slope. there is some evidence for a xeak in the slope at οἱ=0.37. primarily in the IHTorologiua field.," Although the spectrum is consistent with a flat slope, there is some evidence for a break in the slope at $d =
+0.3\dg$, primarily in the Horologium field."
+ For huger aneular scales. the SF i$ approximately Hat iu the Auriga region. with a tentative increase at tle areest lags. aud even decreasing in Horologiuu.," For larger angular scales, the SF is approximately flat in the Auriga region, with a tentative increase at the largest lags, and even decreasing in Horologium."
+ We can estimate he magnitude of the contribution of unreliably determined RASS bv reevaluating the SE or the complete erid of RAL values. iustead of ouly the reliable RASS. This estimate is important for a discussion of the power spectra of RAL. which were evaluated over he complete dataset. including wnreliably determined RALs. The structure function using the complete dataset is shown in the left panel of Fie.," We can estimate the magnitude of the contribution of unreliably determined $RM$ s by reevaluating the SF for the complete grid of $RM$ values, instead of only the reliable $RM$ s. This estimate is important for a discussion of the power spectra of $RM$, which were evaluated over the complete dataset, including unreliably determined $RM$ s. The structure function using the complete dataset is shown in the left panel of Fig."
+ 13 as a dotted line., \ref{f9:sf_rm} as a dotted line.
+" The structure function clearly has a lower amplitude if the unreliable RAL determinations are removed. but the slope of the structure fiction remains approximately the salle,"," The structure function clearly has a lower amplitude if the unreliable $RM$ determinations are removed, but the slope of the structure function remains approximately the same."
+ We compute the structure function of polarized intensity. for both the complete erid of beams. and for those deans selected. to have high. P.," We compute the structure function of polarized intensity, for both the complete grid of beams, and for those beams selected to have high $P$ ."
+ In Fie., In Fig.
+ 11 we show structure functions of P in the Auriga region for 5 frequencies.," \ref{f9:sf_p} we show structure functions of $P$ in the Auriga region for 5 frequencies,"
+ y-ray photons from the Crab pulsar come during GPs.,$\gamma$ -ray photons from the Crab pulsar come during GPs.
+" For our choice of windows, Finax ranged from 60 Fi, (for on-pulse phase window) to 150 F4, (for IP window)."," For our choice of windows, $F_{\mathrm{max}}$ ranged from 60 $F_{\mathrm{av}}$ (for on-pulse phase window) to 150 $F_{\mathrm{av}}$ (for IP window)."
+ The likelihood p(N|F= was calculated by running simulations with different pulsedFo) flux Fo and computing the fraction of runs with a number of matches N., The likelihood $p(N | F = F_{0})$ was calculated by running simulations with different pulsed flux $F_0$ and computing the fraction of runs with a number of matches $N$.
+" The grid of trial flux values, in units of F4,, was as follows: from 0 to 1 with the step of 0.25, from 1 to 20 with the step of 1 or 0.5, and then from 20 to 30 with the step of 5."," The grid of trial flux values, in units of $F_{\mathrm{av}}$, was as follows: from 0 to 1 with the step of 0.25, from 1 to 20 with the step of 1 or 0.5, and then from 20 to 30 with the step of 5."
+ For both choices of correlation window the probability density went down to 0 before 30Εαν.," For both choices of correlation window the probability density went down to 0 before $30\,F_{\mathrm{av}}$."
+ Here we implicitly assumed that a higher flux outside selected windows does not influence the correlation within windows., Here we implicitly assumed that a higher flux outside selected windows does not influence the correlation within windows.
+" Since the number of simulation runs for each trial Fo is finite, it leads to an uncertainty in estimating the likelihood."," Since the number of simulation runs for each trial $F_0$ is finite, it leads to an uncertainty in estimating the likelihood."
+ We estimated the statistical errors from the simulation using the following method., We estimated the statistical errors from the simulation using the following method.
+" Suppose that for some value of Fo we have run n simulations with y successes cases where the number of matches in the simulation (i.e.equals the one obtained for real data, N)."," Suppose that for some value of $F_0$ we have run $n$ simulations with $y$ successes (i.e. cases where the number of matches in the simulation equals the one obtained for real data, $N$ )."
+ Then y/n defines the estimate of probability of success p. which is also the likelihood density p(N|F=Fo).," Then $y/n$ defines the estimate of probability of success $p$ , which is also the likelihood density $p(N | F = F_{0})$."
+" More precisely, p|y has a Beta distribution, with mean (y+1)/(n1) and variance o?="," More precisely, $p|y$ has a Beta distribution, with mean $(y+1)/(n+1)$ and variance $\sigma_p^2=\frac{(y+1)(n-y+1)}{(n+3)(n+2)^2}$."
+ We adopted c; as an error of p due to limited c.numbers of simulations performed., We adopted $\sigma_p$ as an error of $p$ due to limited numbers of simulations performed.
+" However, there is another major source of uncertainty connected to the fact that we record number of photons around the GPs."," However, there is another major source of uncertainty connected to the fact that we record number of photons around the GPs."
+" Since photon detection is very well described as a Poisson process (?),, the error on detecting N photons in a certain window around GPs will be VN."," Since photon detection is very well described as a Poisson process \citep{ramanamurthy1998}, the error on detecting $N$ photons in a certain window around GPs will be $\sqrt{N}$."
+" In our case, for all windows we had N=1, so to estimate the true value of the likelihood we should take into consideration also the likelihood curves for N—0 and N—2."," In our case, for all windows we had $N=1$, so to estimate the true value of the likelihood we should take into consideration also the likelihood curves for $N=0$ and $N=2$."
+" These estimates for both windows are plotted in Fig. 6,,"," These estimates for both windows are plotted in Fig. \ref{fig:correlation-bayes},"
+ left., left.
+" The shaded region around each curve corresponds to co,, calculated by the above formula."," The shaded region around each curve corresponds to $\pm \sigma_{p}$, calculated by the above formula."
+" For both correlation windows N=1 likelihood curves have maximum around =1, which means that, most probably, pulsed y-ray Fo/F,, flux does not change during GPs (no correlation), or changes no more than few times (weak correlation)."," For both correlation windows $N=1$ likelihood curves have maximum around $F_0/F_{\mathrm{av}}=1$, which means that, most probably, pulsed $\gamma$ -ray flux does not change during GPs (no correlation), or changes no more than few times (weak correlation)."
+" With our limited data set we cannot say anything more about the exact value of y-ray flux during GPs, but we can place upper limits on it."," With our limited data set we cannot say anything more about the exact value of $\gamma$ -ray flux during GPs, but we can place upper limits on it."
+" On the grid of simulated fluxes ΕΙ, one can convert the continuous formula for posterior probability density (eq. 2))"," On the grid of simulated fluxes $F_i$, one can convert the continuous formula for posterior probability density (eq. \ref{eq:bayes}) )"
+ into a discrete one for the probability that pulsed flux around GPs is than Fo: where p;=p(N|F , into a discrete one for the probability that pulsed flux around GPs is than $F_0$: where $p_i \equiv p(N | F = F_i)$ .
+"To estimate errors in F;).P(F<Fo|N), we assumed that our uncertainty in p; due to a limited number of trials is much larger than the error from calculating the integral as a sum."," To estimate errors in $P(F \leq F_0 | N)$, we assumed that our uncertainty in $p_i$ due to a limited number of trials is much larger than the error from calculating the integral as a sum."
+" As one can see in Fig. 6,,"," As one can see in Fig. \ref{fig:correlation-bayes},"
+" left, this simplification is reasonable."," left, this simplification is reasonable."
+" Assuming all are independent, the uncertainty in P(F<Fo|N) is op,determined by standard error propagation."," Assuming all $\sigma_{p_i}$ are independent, the uncertainty in $P(F \leq F_0 | N)$ is determined by standard error propagation."
+ In Fig., In Fig.
+ 6 we show the resulting probabilities that the y-ray flux from the Crab pulsar during GPs does not exceed a given number of times the mean flux reported by ?.., \ref{fig:correlation-bayes} we show the resulting probabilities that the $\gamma$ -ray flux from the Crab pulsar during GPs does not exceed a given number of times the mean flux reported by \cite{FermiCrab2010}.
+" Errors due to the limited number of simulation runs are plotted as errorbars, whereas those due to a discrete number of GP/photon matches are given by separate posterior probability curves for N=0 and N=2."," Errors due to the limited number of simulation runs are plotted as errorbars, whereas those due to a discrete number of GP/photon matches are given by separate posterior probability curves for $N=0$ and $N=2$."
+" Obviously, the smaller the correlation window for a fixed observation timespan and the smaller the number of GPs in the sample, the larger Finax, and the broader is the resulting posterior probability density."," Obviously, the smaller the correlation window for a fixed observation timespan and the smaller the number of GPs in the sample, the larger $F_{\mathrm{max}}$, and the broader is the resulting posterior probability density."
+" That is why for our data set we could obtain the posterior probability densities only forcorrelation windows which included the IP,because most of GPs come within this phase range."," That is why for our data set we could obtain the posterior probability densities only forcorrelation windows which included the IP,because most of GPs come within this phase range."
+" For the main pulse GPs, p(F= is very broad, having almost the sameprobability Fo|N)density up"," For the main pulse GPs, $p(F=F_0|N)$ is very broad, having almost the sameprobability density up"
+As it is always compatible with zero and as it yields no improvement of the fits. no global line profile shift was applied. except for the iron lines near wwhere the improvement is significant with a shift of only [542kms.,"As it is always compatible with zero and as it yields no improvement of the fits, no global line profile shift was applied, except for the iron lines near where the improvement is significant with a shift of only $\pm$."
+. Wavelength shifts between instruments/orders were envisaged. but they yielded erratic results generally without improvement of the fit. sometimes large values (incompatible with wavelength calibration reports) for example for the tron line atISA.. values compatible with zero within 3 sigma. and values for one instrument/order compatible with those of another instrument/order within 3 sigma.," Wavelength shifts between instruments/orders were envisaged, but they yielded erratic results generally without improvement of the fit, sometimes large values (incompatible with wavelength calibration reports) for example for the iron line at, values compatible with zero within 3 sigma, and values for one instrument/order compatible with those of another instrument/order within 3 sigma."
+ Cross-correlation also suggests null shifts between exposures and between different instrument/order combinations of the same exposure., Cross-correlation also suggests null shifts between exposures and between different instrument/order combinations of the same exposure.
+ Due to the broadness of its Imes. iis indeed not the best source to find such shifts: to do so. the ccalibration team uses sources with unresolved X-ray lines.," Due to the broadness of its lines, is indeed not the best source to find such shifts: to do so, the calibration team uses sources with unresolved X-ray lines."
+ As shown in the right panel of Fig. 3..," As shown in the right panel of Fig. \ref{veloc},"
+ the optical depth varies with wavelength., the optical depth varies with wavelength.
+ Our fitting thus confirms the Cohenetal.(2010) preliminary results by extending them to longer wavelengths and by decreasing the noise on most lines (except the bluest ones)., Our fitting thus confirms the \citet{coh10} preliminary results by extending them to longer wavelengths and by decreasing the noise on most lines (except the bluest ones).
+ The wind of ttherefore is unlikely to be composed of clumps which are fully opaque at all wavelengths. as had been suggested for some porosity models.," The wind of therefore is unlikely to be composed of clumps which are fully opaque at all wavelengths, as had been suggested for some porosity models."
+ Indeed. optically thick clumps should produce a grey opacity that would solely be determined by the geometry of the clumps.," Indeed, optically thick clumps should produce a grey opacity that would solely be determined by the geometry of the clumps."
+ In principle. these optical depth variations may be used to constrain the mass-loss rate. for a given star-wind model.," In principle, these optical depth variations may be used to constrain the mass-loss rate, for a given star+wind model."
+ This was attempted in Cohenet (2010).. and they found that a reduced mass-loss rate with non- abundances provided the best-fit results.," This was attempted in \citet{coh10}, and they found that a reduced mass-loss rate with non-solar abundances provided the best-fit results."
+ For comparison purposes. the same theoretical opacity is shown in Fig.," For comparison purposes, the same theoretical opacity is shown in Fig."
+ 3. with a dotted line: below20À.. it indeed provides a good fit. but for larger wavelengths. the agreement is less good because of the nitrogen edge at -- an edge which appears very strong as the nitrogen abundance is enhanced forPuppis.," \ref{veloc} with a dotted line: below, it indeed provides a good fit, but for larger wavelengths, the agreement is less good because of the nitrogen edge at - an edge which appears very strong as the nitrogen abundance is enhanced for."
+. This discrepancy may be solved by adapting the value of the abundances and mass-loss rate ofPuppis.. but this is beyond the scope of this paper.," This discrepancy may be solved by adapting the value of the abundances and mass-loss rate of, but this is beyond the scope of this paper."
+ Clearly. additional modelling is needed to reproduce the full behaviour of im X-rays (and this will be done in Paper HD.," Clearly, additional modelling is needed to reproduce the full behaviour of in X-rays (and this will be done in Paper III)."
+ In contrast. the onset radius appears remarkably stable and confined in the R. range. which is quite usual in massive stars (Güdel&Nazé.2009.andreferencestherein)..," In contrast, the onset radius appears remarkably stable and confined in the $R_*$ range, which is quite usual in massive stars \citep[and references therein]{gud09}."
+ The sole exceptions are the Nitrogen triplet and Ένα lines. but it would not be surprising that. as these lines have emissivities that peak at rather low temperatures. it simply indicates a formation further out in the wind.," The sole exceptions are the Nitrogen triplet and $\alpha$ lines, but it would not be surprising that, as these lines have emissivities that peak at rather low temperatures, it simply indicates a formation further out in the wind."
+ In the past decade. about IMs of data were obtained by oonPuppis.," In the past decade, about 1Ms of data were obtained by on."
+. Of these. about is strongly affected by flares. reducing the useful exposures to 579ks for EPIC-MOS. 477ks for EPIC-pn. and 751ks for RGS.," Of these, about is strongly affected by flares, reducing the useful exposures to 579ks for EPIC-MOS, 477ks for EPIC-pn, and 751ks for RGS."
+ A variety of modes was used. the most reliable being the SW+Thick Filter mode: the use of the medium filter vielded piled-up data. while the data taken in LW+Thick Filter mode do not appear significantly different from those obtained with SW-Thick. Filter mode.," A variety of modes was used, the most reliable being the SW+Thick Filter mode; the use of the medium filter yielded piled-up data, while the data taken in LW+Thick Filter mode do not appear significantly different from those obtained with SW+Thick Filter mode."
+ Attention was paid to this problem. notably by choosing similar extraction regions for all datasets and using only single events.," Attention was paid to this problem, notably by choosing similar extraction regions for all datasets and using only single events."
+ Broad-band EPIC data taken with the Thick filter were analyzed using absorbed optically-thin thermal emissions., Broad-band EPIC data taken with the Thick filter were analyzed using absorbed optically-thin thermal emissions.
+ Four temperatures are needed to reproduce in a reasonable way the data., Four temperatures are needed to reproduce in a reasonable way the data.
+ Note that the fits are not formally acceptable. but that (1) instruments do not always agree with one another and (2) the reduced noise amplifies the limitations due to the imperfect atomic parameters and standard line profiles.," Note that the fits are not formally acceptable, but that (1) instruments do not always agree with one another and (2) the reduced noise amplifies the limitations due to the imperfect atomic parameters and standard line profiles."
+ Nevertheless. the EPIC spectra appear remarkably stable over the decade of observations. with only dispersions around the average fluxes.," Nevertheless, the EPIC spectra appear remarkably stable over the decade of observations, with only dispersions around the average fluxes."
+ A detailed variability study. based on lightcurves. will be presented in Paper II.," A detailed variability study, based on lightcurves, will be presented in Paper II."
+ The combined high-resolution RGS spectrum confirms that the X-ray line profiles vary with wavelength., The combined high-resolution RGS spectrum confirms that the X-ray line profiles vary with wavelength.
+ Fitting individual line profiles using a wind model yields similar onset radius for the X-ray emission. but wind continuum opacities depending on wavelength.," Fitting individual line profiles using a wind model yields similar onset radius for the X-ray emission, but wind continuum opacities depending on wavelength."
+ This is simply due to the fact that the cool absorbing clumps in the wind are not fully optically thick at all wavelengths. though further modelling is needed in order to adequately reproduce the opacity variations.," This is simply due to the fact that the cool absorbing clumps in the wind are not fully optically thick at all wavelengths, though further modelling is needed in order to adequately reproduce the opacity variations."
+they have similar values of the dimensionless spin paraiicter a,they have similar values of the dimensionless spin parameter $a_*$.
+ Schuittinan Bertschinger (2001) )orforiued rav tracing caleulations in the Ferr netric for cluitting obs orbiting a black hole a a radius where here is a 3:1 resonance between the azinuthal (Q) and radial coordinate frequencies., Schnittman Bertschinger (2004) performed ray tracing calculations in the Kerr metric for emitting blobs orbiting a black hole at a radius where there is a 3:1 resonance between the azimuthal $\Omega$ ) and radial coordinate frequencies.
+ It was shown that general relativistic effects nav impart à QPO at O and a beat oscillation at 20 with a relative strength that depends ou the angular width of the cmitting blob of uaterial.," It was shown that general relativistic effects may impart a QPO at $\Omega$ and a beat oscillation at $\slantfrac{2}{3}\, \Omega$ with a relative strength that depends on the angular width of the emitting blob of material."
+ This interpretation. when combined with accurate lass nieasurenieuts. vields values for the diueusiouless spin parameter: 0.10<a.<0.55 or GRO J1655-10 aud 0.8<a.0.6 for NTE J1550-56[ C(Renillard et 22002).," This interpretation, when combined with accurate mass measurements, yields values for the dimensionless spin parameter: $0.40 \le a_*\le 0.55$ for GRO J1655-40 and $0.3 \le a_* \le 0.6$ for XTE J1550-564 (Remillard et 2002)."
+ This illustrates he potential diagnostic power of fast QPOs. if we can specifv the correct oscillation mechauisi.," This illustrates the potential diagnostic power of fast QPOs, if we can specify the correct oscillation mechanism."
+ Furthermore. measurements of black hole iiass aud OPO frequency allow colmparisous of spin differences between black hole subclasses as distinguished. for exiiuple. by relativistic jets or binary period.," Furthermore, measurements of black hole mass and QPO frequency allow comparisons of spin differences between black hole subclasses as distinguished, for example, by relativistic jets or binary period."
+ With our measurement of the orbital period and optical mass function of J1650. we have takeu he first steps need to fully exploit the potential of this system.," With our measurement of the orbital period and optical mass function of J1650, we have taken the first steps need to fully exploit the potential of this system."
+ If the QPO at ν=25045 Iz observed bv Woman et ((2003b) represcuts he 21:9 oscillation. then the predicted mass of he black hole would be about 7.5M... if 31650 ies on the same vo:AM curve on which CRO J1655-10. NTE J1550-561. aud CBS 1915|105 lie HE the black holes all have simular values of he dineusiouless spin paraneter «.).," If the QPO at $\nu=250\pm 5$ Hz observed by Homan et (2003b) represents the $2\nu_0$ oscillation, then the predicted mass of the black hole would be about $7.5\,M_{\odot}$, if J1650 lies on the same $\nu:M$ curve on which GRO J1655-40, XTE J1550-564, and GRS 1915+105 lie if the black holes all have similar values of the dimensionless spin parameter $a_*$ )."
+ Although here are substautial uucertaiuties iu both our neasured value for the optical mass function auc in our inclination estimate. it appears that the lack hole 11ass unieght be less than 7.5AZ... which in the context of the orbital resonance mode would imply that the value «. in J1650 is much sunaller than it is for the other three sources AADramowiez et 22001).," Although there are substantial uncertainties in both our measured value for the optical mass function and in our inclination estimate, it appears that the black hole mass might be less than $7.5\,M_{\odot}$, which in the context of the orbital resonance model would imply that the value $a_*$ in J1650 is much smaller than it is for the other three sources Abramowicz et 2004)."
+ Ou the other hand. on the basis of X-ray spectroscopy. Miller OT ((2002) have argued that J1650 contains a maxuual err black hole.," On the other hand, on the basis of X-ray spectroscopy, Miller et (2002) have argued that J1650 contains a maximal Kerr black hole."
+ Thus J1650 may be a challeugiug case for the orbital resonance model of the high fequeney QPOx. although we again poiut out that our ass estimate las large uncertainties.," Thus J1650 may be a challenging case for the orbital resonance model of the high frequency QPOs, although we again point out that our mass estimate has large uncertainties."
+ Recent models for black hole formation FFiver Kalogera 2001) predict a contiuuous and roughly exponcutial distribution of masses for lack holes in binary svsteiis. where the uiuuber of lack holes formed falls off as the mass increases.," Recent models for black hole formation Fryer Kalogera 2001) predict a continuous and roughly exponential distribution of masses for black holes in binary systems, where the number of black holes formed falls off as the mass increases."
+ Depending on their model assunptious. these authors predict the nuuber of black holes with uasses m the range 3.5M. should be roughly wo to three times the ummber of black holes wit[um uassesintherauges 104.," Depending on their model assumptions, these authors predict the number of black holes with masses in the range $3-5\,M_{\odot}$ should be roughly two to three times the number of black holes with masses in the range $8-10\,M_{\odot}$."
+".Dailvn et ((1998 orforiued a statistical analysis of seven black roles with dyvnanucal mass estimates and found strong evidence for a peak in the mass distribution ceutered near 7M. anda ""gap inthe distribution vetween the neutron star masses and the pea- rear TAL..."," Bailyn et (1998) performed a statistical analysis of seven black holes with dynamical mass estimates and found strong evidence for a peak in the mass distribution centered near $7\,M_{\odot}$ and a “gap” in the distribution between the neutron star masses and the peak near $7\,M_{\odot}$."
+ Since that time. the sample of objects ias more than doubled. aud it now appears that he evidence for a poak in the distribution near TAL. is weak (see the references cited in sll ancl Orosz 2003).," Since that time, the sample of objects has more than doubled, and it now appears that the evidence for a peak in the distribution near $7
+\, M_{\odot}$ is weak (see the references cited in 1 and Orosz 2003)."
+ However. black holes with masses inthe 3.SAL. range still seem to be rare. with GRO 10025|32 being the only οσους candidate (Colimo Tarrison 2003).," However, black holes with masses in the $3-5\,M_{\odot}$ range still seem to be rare, with GRO J0422+32 being the only good candidate (Gelino Harrison 2003)."
+ Tn this regard. J1650 is an interesting svstem because of its probable low mass.," In this regard, J1650 is an interesting system because of its probable low mass."
+ Clearly additional data will be required to make more cefinitive statements about the mass of the compact object in J1650., Clearly additional data will be required to make more definitive statements about the mass of the compact object in J1650.
+ It should uot be nuduly difficult to obtain higher quality spectra., It should not be unduly difficult to obtain higher quality spectra.
+ Tn hindsight. we now know that the companion star is a [k-stay. and asx such it should have relatively strong absorption features nearSLTOA.," In hindsight, we now know that the companion star is a K-star, and as such it should have relatively strong absorption features near."
+. This region was not included in the VLT spectra owing to the relatively high resolving power of the 1200R. eria., This region was not included in the VLT spectra owing to the relatively high resolving power of the 1200R grism.
+ It should be easier to measure radial velocities iu spectra with more wavelength coverage (at the expense of resolving powcr)., It should be easier to measure radial velocities in spectra with more wavelength coverage (at the expense of resolving power).
+ Also. spectral observations obtained in good to excelleu seeing conditions (Z0.6 arcsecouds} obviously will have better signal-to-noise than the curren spectra do (the secing varied between 1 aud abou — areseconds).," Also, spectral observations obtained in good to excellent seeing conditions $\lesssim 0.6$ arcseconds) obviously will have better signal-to-noise than the current spectra do (the seeing varied between 1 and about 1.4 arcseconds)."
+ Additional photometry should be obtained. aux it would be helpful if two or more biupasses could be used. since the amount of contaninatiue Πο from the disk should vary with color.," Additional photometry should be obtained, and it would be helpful if two or more bandpasses could be used, since the amount of contaminating light from the disk should vary with color."
+ Using Z-band photometry of NTE J1650- obtained with the σοι Clav. telescope. we have measured a photometric period of," Using $R$ -band photometry of XTE J1650-500 obtained with the 6.5m Clay telescope, we have measured a photometric period of"
+(Biuggelietal.LOSS:Ferguson&Biuggeli1991:GallagherWyse199[).. of (Biuggelietal.1988).. Ixorineudsy.19825).,"\citep{BST88,FB94,GW94}. \citep{BST88}. \citep{LF83,K85},"
+(Biuggelietal.LOSS:Ferguson&Biuggeli1991:GallagherWyse199[).. of (Biuggelietal.1988).. Ixorineudsy.19825)..,"\citep{BST88,FB94,GW94}. \citep{BST88}. \citep{LF83,K85},"
+panels of. 11. we correspondinelv. plot the X-ray.r emission. from. each of. these against the Rossby Number. Na=Pr. where P is the rotation period of the star and 7. is (he convective (turnover limescale. estimated using the relationship from Mim and Demarque (1996).,"panels of 1, we correspondingly plot the X-ray emission from each of these against the Rossby Number, $N_R = P/\tau_c$, where $P$ is the rotation period of the star and $\tau_c$ is the convective turnover timescale, estimated using the relationship from Kim and Demarque (1996)."
+ In these plots. as before. we mark each star by the syaubol C. g. or I/5. based on the classification described earlier.," In these plots, as before, we mark each star by the symbol C, g, or I/S, based on the classification described earlier."
+ One observes that the stars marked C lie in (he super-saturated region of X-ray activity. those marked ο in the saturated region. and those marked I in the region of validity of the rolalion-aclivily relationship.," One observes that the stars marked C lie in the super-saturated region of X-ray activity, those marked g in the saturated region, and those marked I in the region of validity of the rotation-activity relationship."
+ This picture is further reinforced bx ll in BRandich. (2000). which shows that all (he solar-tvpe field stars. expected to lie on the Interface sequence (I) in the new paradigm. obex the standard rotation-activity. relationship.," This picture is further reinforced by 1 in Randich (2000), which shows that all the solar-type field stars, expected to lie on the Interface sequence (I) in the new paradigm, obey the standard rotation-activity relationship."
+ The foregoing observations encourage us to think that the paradigm suggested to account for the rotational evolution of stars mightex also be useful in understandinge the X-ray. emission from stars., The foregoing observations encourage us to think that the paradigm suggested to account for the rotational evolution of stars might also be useful in understanding the X-ray emission from stars.
+ Purely observationally. we know that the A-ray emission V=X(N)A(2/7).," Purely observationally, we know that the X-ray emission $ X = X(N_R) = X(P/\tau_c)$."
+" Now we know from Barnes (2003) Chat P? is not a smoothly. varving function of stellar properties but rather that where P and P; have separate functional dependences on stellar age and color. ancl D, merely represents stars with rotation periods in the gap between the Convective and Interface sequences."," Now we know from Barnes (2003) that $P$ is not a smoothly varying function of stellar properties but rather that where $P_C$ and $P_I$ have separate functional dependences on stellar age and color, and $P_g$ merely represents stars with rotation periods in the gap between the Convective and Interface sequences."
+" Very briefly. ΓΩ... where / is the age. Z, aud I. ave. respectively. (he moments of inertia of the entire star or only the convection zone. and A ds a constant."," Very briefly, $P_I \sim \sqrt{t/I_s}$ but $P_C \sim e^{t/(KI_c)}$ where $t$ is the age, $I_s$ and $I_c$ are, respectively, the moments of inertia of the entire star or only the convection zone, and $K$ is a constant."
+ The origin of these sequences in terms of stellar properties has been explained al length in Barnes (2003)., The origin of these sequences in terms of stellar properties has been explained at length in Barnes (2003).
+ Thus. the fact that the X-ray emission has the particular three-segmented morphology that has been noted by various studies can be traced to the [act Chat. for the three sections.," Thus, the fact that the X-ray emission has the particular three-segmented morphology that has been noted by various studies can be traced to the fact that, for the three sections,"
+with the estimates in Figure 4..,with the estimates in Figure \ref{contour}.
+ While there is no prool for a causal relation between electrons and coronal heating. the energy. budget is compatible with such a scenario for a reasonable range of parameters 1.0.," While there is no proof for a causal relation between high-energy electrons and coronal heating, the energy budget is compatible with such a scenario for a reasonable range of parameters $B, \delta$."
+ We note that. within the framework of our simplified model assumptions. our energy estimate is conservative.," We note that, within the framework of our simplified model assumptions, our energy estimate is conservative."
+ The lifetime 7 of the non-thermal electron population has been adopted at the highest possible value compatible with the radio light curve (1.e.. approximately equal to the shortest decay. Gime scale in the observed radio emission).," The lifetime $\tau$ of the non-thermal electron population has been adopted at the highest possible value compatible with the radio light curve (i.e., approximately equal to the shortest decay time scale in the observed radio emission)."
+ Equation 12. shows that a shorter possible lifetime requires a proportionatelv larger total energy input., Equation \ref{integral} shows that a shorter possible lifetime requires a proportionately larger total energy input.
+ In that case. (he lisht curve variability is controlled by the (time scales of the particle accelerator.," In that case, the light curve variability is controlled by the time scales of the particle accelerator."
+ If electrons get lost after one magnetic-loop crossing (ime. (hen 7 a few seconds for typical loop sizes. aud the rapid replenishment of electrons requires an energy input into tlie svstem up to 3 orders of magnitude higher than estimated above.," If electrons get lost after one magnetic-loop crossing time, then $\tau \approx$ a few seconds for typical loop sizes, and the rapid replenishment of electrons requires an energy input into the system up to 3 orders of magnitude higher than estimated above."
+ We conclude Chat not only does the relative timing between radio and X-ray. emissions support the chromospheric evaporation scenario. but the total energy content in the injected high-energy electrons could easily satisfy or largely exceed the recquirements set bv the observed. X-ray losses.," We conclude that not only does the relative timing between radio and X-ray emissions support the chromospheric evaporation scenario, but the total energy content in the injected high-energy electrons could easily satisfy or largely exceed the requirements set by the observed X-ray losses."
+ The observation described here provides strong support for the chromospheric evaporation scenario in a star that may. maintain a corona considerably. different. [rom the Sun's (e.g.. containing much larger magnetic loops. confining much hotter and perhaps also clenser thermal plasma. magnetic field lines (hat may be arranged in the form of large global dipoles. possibly also between (he companion stars. etc).," The observation described here provides strong support for the chromospheric evaporation scenario in a star that may maintain a corona considerably different from the Sun's (e.g., containing much larger magnetic loops, confining much hotter and perhaps also denser thermal plasma, magnetic field lines that may be arranged in the form of large global dipoles, possibly also between the companion stars, etc)."
+ In retrospect. we find a similar timing between radio and ταν [lare events in some previously published light curves. although the Neupert effect was not. discussed.," In retrospect, we find a similar timing between radio and X-ray flare events in some previously published light curves, although the Neupert effect was not discussed."
+ Most. evidently. radio emission peaking belore the soft X-rays. (hus suggesting the presence of a Neupert effect. can be seen in the examples presented by Vilhuetal.(1933).. Sternetal. (1992).. Brownetal. (1993).. and οἱal. (2001).," Most evidently, radio emission peaking before the soft X-rays, thus suggesting the presence of a Neupert effect, can be seen in the examples presented by \citet{vilhu88}, \citet{stern92}, \citet{brown98}, and \citet{ayres01}."
+. Clearly. differing behavior has been noted as well.," Clearly, differing behavior has been noted as well."
+ First. the Sun shows the Neupert effect most reliably in the class of impulsive flares. whereas of all gradual flares. olten related (o energy release at high coronal altitudes. show a different behavior 1993).," First, the Sun shows the Neupert effect most reliably in the class of impulsive flares, whereas of all gradual flares, often related to energy release at high coronal altitudes, show a different behavior \citep{dennis93}."
+. lt. ds. possible that the connectivity. of magnetic fields between the hieh corona and the chromospheric regions is diflerent in (hese cases. impecding the free [low of electrons aud consequent chromospheric evaporation.," It is possible that the connectivity of magnetic fields between the high corona and the chromospheric regions is different in these cases, impeding the free flow of electrons and consequent chromospheric evaporation."
+ Thermalization of fast electrons could also oceur in the corona already. if the travel distances are long enough and the densities hieh enough., Thermalization of fast electrons could also occur in the corona already if the travel distances are long enough and the densities high enough.
+ Stellar counter-examples of (he Neupert effect include an impulsive optical [lare with following gracual radio emission (vandenOordetal. 1996).. evrosvnchrotron emission," Stellar counter-examples of the Neupert effect include an impulsive optical flare with following gradual radio emission \citep{vdoord96}, , gyrosynchrotron emission"
+Stars with Strónuugren imetallicities higher than about 0.5 dex are supposed to be CN-ricli stars of one of the two populations. since no stars with an iron abundance hieher than that has been found iu the cluster.,"Stars with Strömmgren metallicities higher than about $-0.8$ dex are supposed to be CN-rich stars of one of the two populations, since no stars with an iron abundance higher than that has been found in the cluster."
+ However. many of them are redder than the “main” ROB.," However, many of them are redder than the “main” RGB."
+ Since their (b4) color is not influenced by CN variations. their existence in the CMD can ouly be explaiue by higher iron abundances.," Since their $(b-y)$ color is not influenced by CN variations, their existence in the CMD can only be explained by higher iron abundances."
+ Iu the upper paucl of Fig., In the upper panel of Fig.
+ 6 the nuubers in the metallicity histograun represent about the correct proportion of metal-poor to metal-rich stars. since the stars have been selected. due to a luminosity cut in the CAD (see Fic. 1)).," \ref{hist1} the numbers in the metallicity histogram represent about the correct proportion of metal-poor to metal-rich stars, since the stars have been selected due to a luminosity cut in the CMD (see Fig. \ref{pred}) )."
+ The metal-poor population iu w Cen dominates the metal-rich one by a ratio of about 3:1., The metal-poor population in $\omega$ Cen dominates the metal-rich one by a ratio of about 3:1.
+ The metallicitv. distribution found iu our investigation is qualitatively very similar to that found by Norris et al. (19963) , The metallicity distribution found in our investigation is qualitatively very similar to that found by Norris et al. \cite{norr96}) )
+and Suutzeff Israft (1996)) from their Calcium abundance measurements., and Suntzeff Kraft \cite{sunt96}) ) from their Calcium abundance measurements.
+ Iu Fig., In Fig.
+ 7 we show the moetallicitv distribution of those stars that are ii common in both samples., \ref{feca} we show the metallicity distribution of those stars that are in common in both samples.
+ The Calcium abundances from Suutzeff Ixraft (upper pancl) have been trausformied to iro- abuudances according to the relation eiven in their paper., The Calcium abundances from Suntzeff Kraft (upper panel) have been transformed to iron abundances according to the relation given in their paper.
+ The behaviour of w Cen regarding its relation betwee- Fe abundance and Strónuueren mctallicity is remarkably different frou that of other elobular clusters (open circley. in the lower pauel of Fig. 7..," The behaviour of $\omega$ Cen regarding its relation between Fe abundance and Strömmgren metallicity is remarkably different from that of other globular clusters (open circles in the lower panel of Fig. \ref{feca},"
+" including NGC 633 NGC 3680, NGC 2395, NGC 6397, M22 and M55. taken from ΤΠ]κο 2000))."," including NGC 6334, NGC 3680, NGC 2395, NGC 6397, M22 and M55, taken from Hilker \cite{hilk00a}) )."
+ The straight relation up to -1 dex (ith large scatter towards higher metallicities) is iu striking contrast. for exiunuple. to the situation in M22 (Richter ot al. 1999).," The straight relation up to -1 dex (with large scatter towards higher metallicities) is in striking contrast, for example, to the situation in M22 (Richter et al. \cite{richp}) ),"
+ where there is a considerable scatter at a fixed ou, where there is a considerable scatter at a fixed iron
+all ~1 and the fits are all consistent with each other.,all $\sim 1$ and the fits are all consistent with each other.
+ The process of litting for large-scale. diffuse components and then removing them frou ihe ILEAO map. results in some attenuation of the ACF on angular scales Z107.," The process of fitting for large-scale, diffuse components and then removing them from the $HEAO$ map, results in some attenuation of the ACF on angular scales $\gs 10^{\circ}$."
+ These factors were determined [rom the same Monte Carlo trials that were used to determine the statistical errors and the fits were adjusted: accordingly., These factors were determined from the same Monte Carlo trials that were used to determine the statistical errors and the fits were adjusted accordingly.
+" Even if these factors are ignored. the fit value of 5b, changes by only 3%."," Even if these factors are ignored, the fit value of $b_x$ changes by only $3\%$."
+" To evaluate the level of uneertainty due asvstematic error in (he intensity distribution of ihe XRD. the two ""extreme"" models of Figure 3 were also fit to (he data in the 2.5°<812 interval."," To evaluate the level of uncertainty due a systematic error in the intensity distribution of the XRB, the two “extreme” models of Figure 3 were also fit to the data in the $2.5^{\circ} < \theta < 12^{\circ}$ interval."
+" The biases resulting from these (wo fits are 6,=0.85d:0.13 for the low z model and 5,=1.3640.21 for the high z model.", The biases resulting from these two fits are $b_x = 0.85 \pm 0.13$ for the low z model and $b_x = 1.36 \pm 0.21$ for the high z model.
+ Since these models are somewhat exaggerated. we conclude that thev represent lower and upper limits of svstematic errors due (ο uncertainty in dI/dz.," Since these models are somewhat exaggerated, we conclude that they represent lower and upper limits of systematic errors due to uncertainty in $dI/dz$."
+ A lit to the Veda et al. (, A fit to the Ueda et al. (
+2003) model indicated in Figure 3 results in a similar value of the bias. (hough the precise results depend on how the model is extended to low redshifts (0«z 0.1).,"2003) model indicated in Figure 3 results in a similar value of the bias, though the precise results depend on how the model is extended to low redshifts $0 < z < 0.1$ )."
+ If this model is extended so that the low z behavior is not allowed to fall below that implied by Mivaji et al. (, If this model is extended so that the low z behavior is not allowed to fall below that implied by Miyaji et al. (
+"1994). dI/dz—2.7x10""ergstemο+. then the fit value of 6, becomes 1.1240.17. which is consistent with that implied by our canonical moclel.","1994), $dI/dz = 2.7 \times 10^{-8} erg~s^{-1}cm^{-2}sr^{-1}$, then the fit value of $b_x$ becomes $1.12 \pm 0.17$, which is consistent with that implied by our canonical model."
+ If instead. we use a linear extrapolation to low reclshifts. the bias can be somewhat (c 15%) higher. but the local emissivity of this model would be nearly 26 below that implied by Mivaji et al. (," If instead, we use a linear extrapolation to low redshifts, the bias can be somewhat $\sim 15 \%$ ) higher, but the local emissivity of this model would be nearly $2\sigma$ below that implied by Miyaji et al. ("
+1994).,1994).
+" The ACF on large angular scales is quite sensitive to the contribution of low = sources (roughly half the ACF at 0=4.5° is due to sources wilh z< 0.1). so any error in estimating the low redshilt eutolf in dI/dz could affect the results dramatically,"," The ACF on large angular scales is quite sensitive to the contribution of low $z$ sources (roughly half the ACF at $\theta = 4.5^{\circ}$ is due to sources with $z \ls 0.1$ ), so any error in estimating the low redshift cutoff in $dI/dz$ could affect the results dramatically."
+ By masking sources stronger than 3.x10.Hergs.tem7 we elfectively truncate the intensity distribution at low redshifts., By masking sources stronger than $3 \times 10^{-11}erg~s^{-1}cm^{-2}$ we effectively truncate the intensity distribution at low redshifts.
+ The truncated profiles were determined [rom the [ας eut and the local Iuminosity [function of Steffen et al. (, The truncated profiles were determined from the flux cut and the local luminosity function of Steffen et al. (
+2003).,2003).
+ I£ the value of the flux eut is in error due to. for example. a difference in normalizations of the source catalog used (o make the cuts (Picecinotti οἱ al.," If the value of the flux cut is in error due to, for example, a difference in normalizations of the source catalog used to make the cuts (Piccinotti et al."
+ 1982) and the Stellen et al., 1982) and the Steffen et al.
+ luminosity function. then this would affect the cutoff redshift ancl wotld be translated to an error in the predicted ACF.," luminosity function, then this would affect the cutoff redshift and would be translated to an error in the predicted ACF."
+" In the extreme limit of no Πας eut. ie. no tiruncation of the dZ/dz profile. (he implied X-ray bias is b,=0.90."," In the extreme limit of no flux cut, i.e., no truncation of the $dI/dz$ profile, the implied X-ray bias is $b_x = 0.90$."
+" Ii the other extreme. ie. a [lux cutoff of LxLO!Hergslem?. the implied bias is b,=1.13."," In the other extreme, i.e., a flux cutoff of $1 \times 10^{-11}erg~s^{-1}cm^{-2}$, the implied bias is $b_x = 1.13$."
+ Therelore. it is unlikely that inaceuracy in characterizing the fIux eut is the source of significantsvstematic error.," Therefore, it is unlikely that inaccuracy in characterizing the flux cut is the source of significantsystematic error."
+ Potentiallv more problematic is the redshift distribution of the unresolved. component ol the XRB., Potentially more problematic is the redshift distribution of the unresolved component of the XRB.
+ Worslev et al. (, Worsley et al. (
+2004) found that above 7feV only ~50% of the XRB is resolved: al(hough. (his conclusion must be tempered somewhat by the [act that the brightest sources they considered (in the Lockman Hole) have f[Iuxes of ~LOPergsFem. 7.,"2004) found that above $7~keV$ only $\sim 50\%$ of the XRB is resolved; although, this conclusion must be tempered somewhat by the fact that the brightest sources they considered (in the Lockman Hole) have fluxes of $\sim 10^{-13}~erg~s^{-1}~cm^{-2}$ ."
+ Sources, Sources
+larger fractions of their I in the form of II.,larger fractions of their H in the form of $_{2}$.
+ In addition to carbon and oxveen. interstellar abundances of Mg. Mn. Ni. Cu and Ge have been determined lor the IID 27778 line of sight (Cartledge2002).," In addition to carbon and oxygen, interstellar abundances of Mg, Mn, Ni, Cu and Ge have been determined for the HD 27778 line of sight \citep{car02}."
+. When compared with siehtlines of similar mean density. (hat toward IID 27778 is only unusual in that the gas-phase Ni abundance is ~40% of that along the tvpical interstellar path.," When compared with sightlines of similar mean density, that toward HD 27778 is only unusual in that the gas-phase Ni abundance is $\sim 40$ of that along the typical interstellar path."
+ The molecule CO has a highly unusual abundance toward ID 27778 with N(CO) = 2.5xLOM 7 (Federmanοἱal.1994:JosephetLO8G).," The molecule CO has a highly unusual abundance toward HD 27778 with N(CO) = $2.5 
+\times 10^{16}$ $^{-2}$ \citep{fed94,jos86}."
+. This high CO column density suggests Chat = of the gas-phase C along the sightline is in the form of CO. and that the CO/IIs ratio is ~ 10 times larger than for any of the other sightlines in Table 2.," This high CO column density suggests that $\gtrsim$ of the gas-phase C along the sightline is in the form of CO, and that the $_{2}$ ratio is $\sim$ 10 times larger than for any of the other sightlines in Table 2."
+ The only other sightline [rom the entire sample that approaches this CO abundanuce is (hat toward ILD 24534., The only other sightline from the entire sample that approaches this CO abundance is that toward HD 24534.
+ The measured N(CO) =1.4x10! (Shefferetal.2002).. however. still puts the HD 24534 sightline CO/IIs a [actor of 2.8 below the ID 27778 value and suggests that only ~4% of the gas-phase CC is in the form of CO.," The measured N(CO) $= 1.4 \times 10^{16}$ \citep{she02}, however, still puts the HD 24534 sightline $_{2}$ a factor of 2.8 below the HD 27778 value and suggests that only $\sim 4\%$ of the gas-phase C is in the form of CO."
+ The unusual CO abundances towird IID 27778 are particularly difficult to understand given that no more than ~ oof the material along the sightline can be in a translucent-like component 2002)., The unusual CO abundances toward HD 27778 are particularly difficult to understand given that no more than $\sim$ of the material along the sightline can be in a translucent-like component \citep{rac02}.
+. Other than the interstellar gas-phase C. CO and Ni abundances. there appears to be no indication from observations or chemical modeling (e.g.. Federmanetal. (1994))) that ihe HD 27778 sightline is unusual.," Other than the interstellar gas-phase C, CO and Ni abundances, there appears to be no indication from observations or chemical modeling (e.g., \citet{fed94}) ) that the HD 27778 sightline is unusual."
+ Although we do not put forth an explanation for it. the fact that (vo carbon-based characteristics. C/II and /II». are highly atypical toward LD P7TTS suggests (hat Chev might be linked.," Although we do not put forth an explanation for it, the fact that two carbon-based characteristics, C/H and $_{2}$, are highly atypical toward HD 27778 suggests that they might be linked."
+ Understanding that link should provide information about the conditions under which carbon depletes bevond the dilluse-cloud level., Understanding that link should provide information about the conditions under which carbon depletes beyond the diffuse-cloud level.
+ We thank the referee for his thoughtful comments., We thank the referee for his thoughtful comments.
+ This work was supported by grants from the Space Telescope Science Iistitute to Whitman College and Northwestern University., This work was supported by grants from the Space Telescope Science Institute to Whitman College and Northwestern University.
+ This research has made use of the SIAIBAD database. operated at CDS. Strasbourg. France.," This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
+of accretion does not imply in the lack of circumstellar material. as verilied by André&Montmerle(1904). in a comparison between the near-infrared and Elo classifications of TTs.,"of accretion does not imply in the lack of circumstellar material, as verified by \cite{Andre94} in a comparison between the near-infrared and $\alpha$ classifications of TTs."
+ The contribution of the circumstellar matter to the SIED slope is often used to recognize different categories of voung stellar objects (YSOs) by following an observational Classification basecl on the near-infrared spectral index (Lada&Wilking1984:etal.1989:Anelréοἱal. 1993).," The contribution of the circumstellar matter to the SED slope is often used to recognize different categories of young stellar objects (YSOs) by following an observational classification based on the near-infrared spectral index \citep{Lada84,Wilking89,Andre93}."
+. This classification suggests a scenario for the evolution of YSOs. from Class 0 to Class LH. which is well established for T's.," This classification suggests a scenario for the evolution of YSOs, from Class 0 to Class III, which is well established for TTs."
+ Although the present work is focused on LLAeBes. it is useful to describe the meaning of the dillercnt T's classes in order to illustrate how the spectral classification relates to cise structure and evolutionary sequence.," Although the present work is focused on HAeBes, it is useful to describe the meaning of the different TTs classes in order to illustrate how the spectral classification relates to disc structure and evolutionary sequence."
+ We summarize here the phases of pre-main sequence evolution. owed: Upon the review by Feigelson& Montmerle(1999) and following the schematic illustration by Carkner(1998).," We summarize here the phases of pre-main sequence evolution, based upon the review by \cite{Feigelson99} and following the schematic illustration by \cite{Carkner98}."
+. The earlier. evolutive phases are. represented by Class 0 (infalling protostar) and Class L (evolved: protostar), The earlier evolutive phases are represented by Class 0 (infalling protostar) and Class I (evolved protostar).
+ 30th contain objects that are surrounded by circumstellar shell and show increasing SED in the infrared range., Both contain objects that are surrounded by circumstellar shell and show increasing SED in the infrared range.
+ Most. of he Class HE objects are CPTs. which are expected to have an accretion disce ancl to show Ilattened spectral index.," Most of the Class II objects are CTTs, which are expected to have an accretion disc and to show flattened spectral index."
+ The ater phase is represented by Class LLL in which are included WES and post-E'Fs.," The later phase is represented by Class III, in which are included WTTs and post-TTs."
+ The disc is thin or non-existent: in Class LLL objects. since they have little or no infrared-excess and show steeper slope of the SED.," The disc is thin or non-existent in Class III objects, since they have little or no infrared-excess and show steeper slope of the SED."
+ For ILXeDBes the spectral classification and its relation with an evolutionary sequence are less advanced than the TVs classes., For HAeBes the spectral classification and its relation with an evolutionary sequence are less advanced than the TTs classes.
+ In Sect., In Sect.
+ 1.2. we present the LLAcBes groups and a tentative scenario for their evolution. as proposed by other authors.," \ref{sectgroups} we present the HAeBes groups and a tentative scenario for their evolution, as proposed by other authors."
+ The Hat disc geometry was also. adopted. to model. the protoplanctary dises of ILXeDes. candidates. detected. by, The flat disc geometry was also adopted to model the protoplanetary discs of HAeBes candidates detected by.
+PDS. Sartorietal.(2003). (hereafter. SCLIO3) classified these stars in different groups. based. on the SED shapes and the fraction of circumstellar luminosity. expressed. by fo=ChinLMLoa.," \cite{Sartori03} (hereafter SGH03) classified these stars in different groups, based on the SED shapes and the fraction of circumstellar luminosity, expressed by $f_c = (L_{\rm tot} - L_{\star})/L_{\rm tot}$."
+" Adopting the spectral index n=0.T5logCFioFT)1 Clorres1998).. the stars were classified as follows: the higher infrared excesses (f. το) are found in the 2,>0 group. which contains 44 per cent of the sample with increasing spectrum in the visible to infrared range: the l«3j«0 group has 50 per cent of the stars showing intermediate amounts of circumstellar emission (f.= 10 to 70 1): in the third. group (4,< 1) there is 6 per cent of the stars having the lower values of infrared excess Cf. 10)."," Adopting the spectral index $\beta_1 = 0.75 {\rm log}(F_{12}/F_V)-1$ \citep{Torres98}, the stars were classified as follows: the higher infrared excesses $f_c>70\%$ ) are found in the $\beta_1 > 0$ group, which contains 44 per cent of the sample with increasing spectrum in the visible to infrared range; the $-1 < \beta_1 < 0$ group has 50 per cent of the stars showing intermediate amounts of circumstellar emission $f_c =$ 10 to 70 ); in the third group $\beta_1 < -1$ ) there is 6 per cent of the stars having the lower values of infrared excess $f_c < 10 \%$ )."
+ Phe groups were established. following the idea of an evolutionary scenario for LLAcBes proposed by. Malfait (1998).. similar to the TPs classes.," The groups were established following the idea of an evolutionary scenario for HAeBes proposed by \cite{Malfait98}, similar to the TTs classes."
+ In this scenario. values of 3)>0.7 correspond to embedded objects. like R ο πα «— Lrepresents the stars showing more evolved disces. like Vega.," In this scenario, values of $\beta_1 > 0.7$ correspond to embedded objects, like R CrA; and $\beta_1 < -1$ represents the stars showing more evolved discs, like Vega."
+ The intermediary group has objects with SED similar to several known LLYeDes., The intermediary group has objects with SED similar to several known HAeBes.
+ This classification is almost the same as the ILXeDe eroups presented by Lillenbrandetal.(1992).. based on the slope of the infrared. continuum.," This classification is almost the same as the HAeBe groups presented by \cite{Hillenbrand92}, based on the slope of the infrared continuum."
+ It is also comparable to the groups proposed by Moeeusetal.(2001). that studied a sample of 14 isolated. Lerbig stars., It is also comparable to the groups proposed by \cite{Meeus01} that studied a sample of 14 isolated Herbig stars.
+ Their two major groups suggest different circumstellar ecometries: the Croup ] sources are associated with an optically thin. Hared region surrounding an optically thick disc. while Group LL sources have a Lat clise (Alecusetal.2002).," Their two major groups suggest different circumstellar geometries: the Group I sources are associated with an optically thin, flared region surrounding an optically thick disc, while Group II sources have a flat disc \citep{Meeus02}."
+. One of the dillerences of this comparison is related to the SCIIO3 group of embedded objects., One of the differences of this comparison is related to the SGH03 group of embedded objects.
+ Since they cannot. be considered: isolated: LLerbig stars they were not classified by Aleeus et al., Since they cannot be considered isolated Herbig stars they were not classified by Meeus et al.
+.. Another difference is the subdivision based. upon the presence (a) or the absence (b) of the 10 fina silicate feature. which is not considered by SCLIIO3.," Another difference is the subdivision based upon the presence (a) or the absence (b) of the 10 $\mu$ m silicate feature, which is not considered by SGH03."
+ Table 1 summarizes (in the last columns) the comparison of the groups above described., Table \ref{tabstars} summarizes (in the last columns) the comparison of the groups above described.
+ SCOGLIOS. estimated. the circumstellar luminosity of thePDS WAcBes adopting the CGLIO2 flat. disc model., SGH03 estimated the circumstellar luminosity of the HAeBes adopting the GH02 flat disc model.
+ They verified that the fitting-qualitv was gof2OA (a hac data fitting) for more than 30 per cent of the objects., They verified that the fitting-quality was $gof > 0.4$ (a bad data fitting) for more than 30 per cent of the objects.
+ This suggests that different disc geometry. should be more appropriate for an important part of the sample., This suggests that different disc geometry should be more appropriate for an important part of the sample.
+ We decided to adopt a IHared configuration. bv following the Dullemondetal.(2001). (hereafter. DDNOL) moclel for a passively irracliatecl circumstellar cise with an inner hole.," We decided to adopt a flared configuration, by following the \cite{DDN2001} (hereafter DDN01) model for a passively irradiated circumstellar disc with an inner hole."
+ Another improvement on our previous work is the method to optimise the parameters estimation based on genetic algorithms (CX)., Another improvement on our previous work is the method to optimise the parameters estimation based on genetic algorithms (GA).
+ This automatic us minimisation procedure was developed to increase the efficiency in the process of the SED fitting., This automatic $\chi^{2}$ minimisation procedure was developed to increase the efficiency in the process of the SED fitting.
+ Phe main goals of the present work are: (i) to describe the implementation of the CiX method for the Πάνο. disc model: (ii) to discuss the derived disc parameters in a variety of SED shapes., The main goals of the present work are: (i) to describe the implementation of the GA method for the flared disc model; (ii) to discuss the derived disc parameters in a variety of SED shapes.
+ Once the GA method has been confirmed to be a valuable tool for SED fittings. the scope of the second goal of this paper is to discuss the value of dilferent spectral classifications in determining disc structures: the relevance of Dat Dared. gcomoetries: and on cdegencracy of the parameters derived from SED mocdelling.," Once the GA method has been confirmed to be a valuable tool for SED fittings, the scope of the second goal of this paper is to discuss the value of different spectral classifications in determining disc structures; the relevance of flat flared geometries; and on degeneracy of the parameters derived from SED modelling."
+ We organized the structure of the paper in the following sequence., We organized the structure of the paper in the following sequence.
+ La Sect., In Sect.
+ 2. we present the adaptation. of the DDNOI model. for which a new code was developed.," \ref{sectmodel} we present the adaptation of the DDN01 model, for which a new code was developed."
+ A test of the rewritten routines was made by applying them to AD Aur., A test of the rewritten routines was made by applying them to AB Aur.
+ Sect., Sect.
+ 3. deseribes the CoA method of parameters estimation and its implementation., \ref{sectga} describes the GA method of parameters estimation and its implementation.
+ Time efficiency and error estimative provided by the CX method. are. discussed. in a comparison with our previous calculation procedure., Time efficiency and error estimative provided by the GA method are discussed in a comparison with our previous calculation procedure.
+ In order to illustrate the application of the GA method to fit dillerent. SED shapes. using the DDNOI model. we present in Sect.," In order to illustrate the application of the GA method to fit different SED shapes, using the DDN01 model, we present in Sect."
+ 4. the results obtained for some of thePLS stars., \ref{secttest} the results obtained for some of the stars.
+" The sample was selected according cdillerent. 3, ranges to represent dillerent groups of ILXeDe stars.", The sample was selected according different $\beta_1$ ranges to represent different groups of HAeBe stars.
+ We choose to study only a few stars in this first. presentation of the GA method just to verily if the adopted dise geometry. is suited to fit different SED shapes., We choose to study only a few stars in this first presentation of the GA method just to verify if the adopted disc geometry is suited to fit different SED shapes.
+ In a forthcoming paper we intend to apply the CX method to evaluate the circumstellar structure of thePOS stars that probably. have Iared disces., In a forthcoming paper we intend to apply the GA method to evaluate the circumstellar structure of the stars that probably have flared discs.
+ Considering the dillerent levels of infrared excess shown by the selected sample. we discuss in Sect.," Considering the different levels of infrared excess shown by the selected sample, we discuss in Sect."
+ 5. the cireumstellar luminosity compared with the decreasing slope of the SED. which has been interpreted as a sequence for the. disc evolution.," \ref{sectdiscus} the circumstellar luminosity compared with the decreasing slope of the SED, which has been interpreted as a sequence for the disc evolution."
+ Ehe quality of the SED fitting is used to compare different morphologies assumed. in the Ilared. and. the fat, The quality of the SED fitting is used to compare different morphologies assumed in the flared and the flat
+"(90 um) and (140 um) fluxes, assuming a spectral shape of I,οςv?B,(T4), where Τα is the dust temperature, 6 is the emissivity index, and B,(T) is the Planck function.","(90 $\micron$ ) and (140 $\micron$ ) fluxes, assuming a spectral shape of $I_\nu\propto \nu^\beta B_\nu (T_\mathrm{d})$, where $T_\mathrm{d}$ is the dust temperature, $\beta$ is the emissivity index, and $B_\nu (T)$ is the Planck function."
+ The emissivity index ( is chosen to be 2 in this paper unless otherwise stated., The emissivity index $\beta$ is chosen to be 2 in this paper unless otherwise stated.
+" We apply Ty=27 K and 19 K for andWIDE-L,, respectively; the former (latter) temperature is derived from the total fluxes of 881 at N660 and and WIDE-L)) (see Section 3.2))."," We apply $T_\mathrm{d}=27$ K and 19 K for and, respectively; the former (latter) temperature is derived from the total fluxes of 81 at 60 and and ) (see Section \ref{subsec:global}) )."
+" Consequently, the correction factor is 0.92 and 0.94 for the and fluxes, respectively, and the flux obtained by the Slow Scan Tool is divided by these factors."," Consequently, the correction factor is 0.92 and 0.94 for the and fluxes, respectively, and the flux obtained by the Slow Scan Tool is divided by these factors."
+ The uncertainty caused by the colour correction (5 per cent for and 1 per cent for for the temperature range observed in the 881 disc; Section ??)) is smaller than the errors caused by the background fluctuation., The uncertainty caused by the colour correction (5 per cent for and 1 per cent for for the temperature range observed in the 81 disc; Section \ref{subsec:Tmap}) ) is smaller than the errors caused by the background fluctuation.
+" For the narrow band N660 (65 ium) we do not apply the correction, since the colour correction only changes the flux by less than 3 per cent."," For the narrow band 60 (65 $\micron$ ) we do not apply the correction, since the colour correction only changes the flux by less than 3 per cent."
+ Bright stripes along the scan direction caused by glitches or non-uniform detector sensitivity are observed in all the three bands., Bright stripes along the scan direction caused by glitches or non-uniform detector sensitivity are observed in all the three bands.
+" In order to eliminate these structures, the background levels are estimated for each line along the scan direction by averaging the intensities in the two sections before and after the scan of the 881main body with a separation of about 18' and a total length of about 20', and then subtracted."," In order to eliminate these structures, the background levels are estimated for each line along the scan direction by averaging the intensities in the two sections before and after the scan of the 81main body with a separation of about $18'$ and a total length of about $20'$, and then subtracted."
+" The RMSs of the background are 1.3, 0.7, 2.9 MJy sr for the 65, 90, and 140 um bands, respectively."," The RMSs of the background are 1.3, 0.7, 2.9 MJy $^{-1}$ for the 65, 90, and 140 $\micron$ bands, respectively."
+ These values are estimated before eliminating the stripes to show the original uncertainty including the non-uniform sensitivity and the time variability of the detector response., These values are estimated before eliminating the stripes to show the original uncertainty including the non-uniform sensitivity and the time variability of the detector response.
+" We matched the positions of the images taken by two detectors, SW (N660 and WIDE-S)) and LW (WIDE-L)) The uncertainty in the relative position between the detectors is well below the grid size."," We matched the positions of the images taken by two detectors, SW 60 and ) and LW ) The uncertainty in the relative position between the detectors is well below the grid size."
+" The intensity (surface brightness) at a wavelength A (frequency V= c/A, where c is the light speed) in each grid is denoted as I,()."," The intensity (surface brightness) at a wavelength $\lambda$ (frequency $\nu\equiv c/\lambda$ , where $c$ is the light speed) in each grid is denoted as $I_\nu (\lambda )$."
+" The intensity ratio, I,(41)/(A2), is called A1—A2 colour in this paper."," The intensity ratio, $I_\nu (\lambda_1)/I_\nu (\lambda_2)$, is called $\lambda_1-\lambda_2$ colour in this paper."
+" We derive the dust temperatureI, 74 in each grid as follows.", We derive the dust temperature $T_\mathrm{d}$ in each grid as follows.
+ It is convenient to convert τ(λ)=Av? to a commonly used indicator of dust optical depth., It is convenient to convert $\tau (\lambda)=A\nu^2$ to a commonly used indicator of dust optical depth.
+ We choose Ay (the extinction in V band in units of magnitude) for such an indicator., We choose $A_V$ (the extinction in $V$ band in units of magnitude) for such an indicator.
+" We adopt the Galactic extinction curve for the conversion from τ(λ) to Ay as derived by Weingartner&Draine(2001) for Ry= 3.1: Αν=C(A)r(A), where the factor C(A) is 8.3x10? for umand1.9x10? for um."," We adopt the Galactic extinction curve for the conversion from $\tau (\lambda )$ to $A_V$ as derived by \citet{weingartner01} for $R_V=3.1$ : $A_V=\mathcal{C}(\lambda )\tau (\lambda )$, where the factor $\mathcal{C}(\lambda )$ is $8.3\times 10^2$ for $\micron$ and $1.9\times 10^3$ for $\micron$ ."
+" AlthoughwederiveAy by using the umvalues, f weobtainthesamevalueorAv. also from the umvaluesbecauseC(A)οςA? holds within 10 per cent between 90 and um."," Although we derive $A_V$ by using the $\micron$values, we obtain the same value for $A_V$ also from the $\micron$ values because $\mathcal{C}(\lambda )\propto\lambda^{2}$ holds within 10 per cent between 90 and $\micron$ ."
+eyisocles Of much. more noderae activity.,episodes of much more moderate activity.
+ None of them slows any evidence for quiesceut pliases onger than 408 vr., None of them shows any evidence for quiescent phases longer than $^8$ yr.
+ If we combine the results o the SF history of the late-ype dwarls studied by our grouy with ose obtained by other groups (see e.g. Grebe 1998 for a colpreheusive review of Local Croup vlaxies and Schulte-Ladbeck. his volume. for BCDs οιtsicle the Local Group) it is apparett that ese gPalaxies. with very lew exceptious. have |ad qualitatively similar SF histories.," If we combine the results on the SF history of the late-type dwarfs studied by our group with those obtained by other groups (see e.g. Grebel 1998 for a comprehensive review of Local Group galaxies and Schulte-Ladbeck, this volume, for BCDs outside the Local Group) it is apparent that these galaxies, with very few exceptions, have had qualitatively similar SF histories."
+ Fie.5 shows e SF histories derived iu 5 gaaxies: NCC‘1569 is the only one which appears to follow the bursting Sheme., \ref{sft} shows the SF histories derived in 5 galaxies: NGC1569 is the only one which appears to follow the bursting scheme.
+ The data available so ar have not allowed vet tο accurately cle‘Ive linearlier activity. but e new HST photometry in tje near infrared yy Aloisi e al. (," The data available so far have not allowed yet to accurately derive its earlier activity, but the new HST photometry in the near infrared by Aloisi et al. ("
+2001) s10uld alow for a lookback ime of. at east. a lew Crs.,"2001) should allow for a lookback time of, at least, a few Gyrs."
+ Il sunmiiarvy: From the above presentation. it is apparent that further studies ou several issues are still required to obtain a reliable scenario for the evolution of late-type dwarfs.," In summary: From the above presentation, it is apparent that further studies on several issues are still required to obtain a reliable scenario for the evolution of late-type dwarfs."
+ Lu particular. the SF histories aud the major characteristics of the gas flows triggered by SN explosions should be examined iu more details.," In particular, the SF histories and the major characteristics of the gas flows triggered by SN explosions should be examined in more details."
+ For this reason I believe that. at present. the correct approach to model the chemical evolution of these systems would be to concentrate ou single representative cases. rather than to moclel the overall features of a large sample of galaxies.," For this reason I believe that, at present, the correct approach to model the chemical evolution of these systems would be to concentrate on single representative cases, rather than to model the overall features of a large sample of galaxies."
+ For the clioseu cases I would:, For the chosen cases I would:
+"can cause a precession rate d»=OSG""μη assuming that it is well constrained.","can cause a precession rate $\dot{\omega}_2=-0.86^\circ/yr$, assuming that it is well constrained."
+ Moreover. we expect our model to explain better the synthetic data generated with 3-bocdy simulations than real data (7-Octantis) where we could have other planets or even stellar variability.," Moreover, we expect our model to explain better the synthetic data generated with 3-body simulations than real data $\nu$ -Octantis) where we could have other planets or even stellar variability."
+ A planet around the primary star in #-Octantis can also cause precession of the main binary's orbit., A planet around the primary star in $\nu$ -Octantis can also cause precession of the main binary's orbit.
+" However. our simulations show that coplanar retrograde planet orbits. as reported in ον, cause slow prograde precession of the main binarys orbit at a rate 0.04 vr."," However, our simulations show that coplanar retrograde planet orbits, as reported in \citet{Eberle&Cuntz2010}, cause slow prograde precession of the main binary's orbit at a rate $0.04^\circ$ /yr."
+ This is also what we expect [rom our quadrupole order theory (Eqs. 25.. 26))," This is also what we expect from our quadrupole order theory (Eqs. \ref{om2dotgen0}, \ref{atheta}) )"
+ although we do not expect it to be accurate at semi-major axis ratio a—0.47., although we do not expect it to be accurate at semi-major axis ratio $\alpha=0.47$.
+ We saw (Sect., We saw (Sect.
+ 23) that in order to have retrograde precession we would. need the planct’s orbit to be inclined more than 45 with respect to the z-Octantis binary., 2.3) that in order to have retrograde precession we would need the planet's orbit to be inclined more than $45^\circ$ with respect to the $\nu$ -Octantis binary.
+ In our numerical integrations we could not find. (although we cid not do an exhaustive search) stable planet orbits at. semi-major axis ratio a=OAT and with such high inclination with respect to the main binary., In our numerical integrations we could not find (although we did not do an exhaustive search) stable planet orbits at semi-major axis ratio $\alpha=0.47$ and with such high inclination with respect to the main binary.
+ We studied the elfect of a binary svstem on a nearby star's motion., We studied the effect of a binary system on a nearby star's motion.
+ This is a complement of our previous work (2) where we assumed that we hacl observations for a fraction of the stars orbit around the binary’s centre of mass.," This is a complement of our previous work \citep{Morais&Correia2008,Morais&Correia2011} where we assumed that we had observations for a fraction of the star's orbit around the binary's centre of mass."
+ Llere. we assumed that we had observations for a few orbits of the star around the binarv's centre of mass.," Here, we assumed that we had observations for a few orbits of the star around the binary's centre of mass."
+ We saw that. in this case. the secular elfect of the binary dominates over the short-term effects.," We saw that, in this case, the secular effect of the binary dominates over the short-term effects."
+ We developed a secular theory which was based on a quacrupole expansion of the Hamiltonian., We developed a secular theory which was based on a quadrupole expansion of the Hamiltonian.
+ This is accurate for hierarchical triple svstems composed of an inner binary. and a star that moves around this inner binary’s centre of mass on à wider orbit which we called the outer binary.," This is accurate for hierarchical triple systems composed of an inner binary, and a star that moves around this inner binary's centre of mass on a wider orbit which we called the outer binary."
+ We derived. an expression. for. the outer binary’s precession rate anc showed that it is) approximately constant., We derived an expression for the outer binary's precession rate and showed that it is approximately constant.
+" ""Therefore. the star's racial velocity can be modeled as a modified. Keplerian racial velocity curve with slowly drifting amplitude."," Therefore, the star's radial velocity can be modeled as a modified Keplerian radial velocity curve with slowly drifting amplitude."
+ We then showed how we can measure the outer binary’s precession rate by fitting a oecessing Ixeplerian. orbit to the radial velocity data., We then showed how we can measure the outer binary's precession rate by fitting a precessing Keplerian orbit to the radial velocity data.
+ We also showed how we can estimate inner binary parameters rom the measured. precession rate., We also showed how we can estimate inner binary parameters from the measured precession rate.
+ We saw that. if we are unaware of the inner binary’s existence and simply fit à non-precessing Ixeplerian orbit to he racial velocity data. à periodogram of the residuals will show peaks at or nearby harmonics of the outer binary's »eriod which can be mistaken as planets.," We saw that, if we are unaware of the inner binary's existence and simply fit a non-precessing Keplerian orbit to the radial velocity data, a periodogram of the residuals will show peaks at or nearby harmonics of the outer binary's period which can be mistaken as planets."
+ However. if we fit a precessine Ixeplerian orbit to the radial velocity data. these," However, if we fit a precessing Keplerian orbit to the radial velocity data, these"
+fraction of ionizing photons Irom their host halos. both due to poorly constrained. physics ab sub-galactie scales and also due to the present lack of spatial resolution.,"fraction of ionizing photons from their host halos, both due to poorly constrained physics at sub-galactic scales and also due to the present lack of spatial resolution."
+ Instead. we have to rely on parameterizing (he escape fraction within an effective UV efficiency factor., Instead we have to rely on parameterizing the escape fraction within an effective UV efficiency factor.
+ llowever. the technique deseribed here when combined with the power of AMIR. within hvdro simulations — can bring us closer to a correct quantitative description of the relevant ealaxv formation processes. both prior to ancl after (he beginning of cosmic reionizalion.," However, the technique described here – when combined with the power of AMR within hydro simulations – can bring us closer to a correct quantitative description of the relevant galaxy formation processes, both prior to and after the beginning of cosmic reionization."
+ Ideally. these models should be combined with the full [requency-dependent. physics of a code such as (ILaardt Madan 2001) to perform exact matching between numerical star formation aud the observed quantities at lower redshifts. such as the UV background or directly measured Iuminosity functions.," Ideally, these models should be combined with the full frequency-dependent physics of a code such as (Haardt Madau 2001) to perform exact matching between numerical star formation and the observed quantities at lower redshifts, such as the UV background or directly measured luminosity functions."
+ Recently obtained spectra of the highest redshift quasars (Becker et al., Recently obtained spectra of the highest redshift quasars (Becker et al.
+ 2001. Djorgovski et al.," 2001, Djorgovski et al."
+ 2001) show tantalizing evidence for the end of the reionizalion process occurring al zo6., 2001) show tantalizing evidence for the end of the reionization process occurring at $z\simeq6$.
+ It mav therefore be possible to obtain data in the near future which directly. probes (he reionizalion process., It may therefore be possible to obtain data in the near future which directly probes the reionization process.
+ The redshift evolution and spatial topology of reionization may be measurable using sullicientlv many lines of sight., The redshift evolution and spatial topology of reionization may be measurable using sufficiently many lines of sight.
+ Comparison with models such as ours max allow connections to be made between the observational diagnostics and physical inputs to the reionization process: IMFEs. SF rates. relative contributions of ACN. sizes of the first collapsing mini-halos. etc.," Comparison with models such as ours may allow connections to be made between the observational diagnostics and physical inputs to the reionization process: IMFs, SF rates, relative contributions of AGN, sizes of the first collapsing mini-halos, etc."
+ A suite of simulations along the lines of those described here will be extvemely helpful in fulfilling Chis goal., A suite of simulations along the lines of those described here will be extremely helpful in fulfilling this goal.
+are depleted. by a factor of 10 since thev are likely to be found in (he form of grains in 11 11 regions (Savage&Sembach1996).,are depleted by a factor of 10 since they are likely to be found in the form of grains in H II regions \citep{savage96}.
+". Accordingly. erains were included in our simulations bv turning on (he ""grains ISAT feature implemented inCLOUDY.. with their abundances scaled down bv a [actor of 1/5. 1/20. or 1/50. in order to be consistent. with the oxveen depletion of themodel."," Accordingly, grains were included in our simulations by turning on the “grains ISM” feature implemented in, with their abundances scaled down by a factor of 1/5, 1/20, or 1/50, in order to be consistent with the oxygen depletion of the."
+ In order to study the effect of varving the nebulas electron density. we ran models [ου two different values of6(H).. 10 or 1007. which are (tvpical values measured in low metallicity systems.," In order to study the effect of varying the nebula's electron density, we ran models for two different values of, 10 or 100, which are typical values measured in low metallicity systems."
+ We also varied the distance [rom the star to the illuminated face of the cloud. (5). adopting either 2 pc. (which is bieger than the radius of a planetary nebula) or 5 pc (which allowed us to cover a range for the ionization parameter comparable (to what is found in the literature).," We also varied the distance from the star to the illuminated face of the cloud $_0$ ), adopting either 2 pc, (which is bigger than the radius of a planetary nebula) or 5 pc (which allowed us to cover a range for the ionization parameter comparable to what is found in the literature)."
+ The above choices resulted in -3.17 < log(U) x -0.42., The above choices resulted in -3.17 $\leq$ log(U) $\leq$ -0.42.
+ The parameters characterizing our models are summarized in Table 3.. which gives O/IILin column (2). the II density in column (3). ry in column (4). log(U) in column (5) and the nebular eeomelrv in column (6).," The parameters characterizing our models are summarized in Table \ref{clouds}, which gives O/H in column (2), the H density in column (3), $_0$ in column (4), log(U) in column (5) and the nebular geometry in column (6)."
+" Note that the latter is plane-parallel if Ar/r,<0.1. athick shellif 0.1<Ar/r,«3. and spherical if Ar/r,>J)3. where Ar is the thickness of the eloud set by the stopping temperature. chosen to be 100 Kx. In order to find parameterizations of and with respect to2).. [vom models.)..).. and need to be determined independent of each other first."," Note that the latter is plane-parallel if $\Delta r/r_o < 0.1$, athick shellif $0.1 < \Delta r/r_o < 3$, and spherical if $\Delta r/r_o\geq 3$, where $\Delta r$ is the thickness of the cloud set by the stopping temperature, chosen to be 100 K. In order to find parameterizations of and with respect to, from models, and need to be determined independent of each other first."
+ This can be accomplished using two different methods., This can be accomplished using two different methods.
+ The first method. corresponding to common practice (Garnett1990:StasitiskaStasinska]lzotov 2003).. is to compute each temperature from: ie. each temperature is an ion and weighted volume mean.," The first method, corresponding to common practice \citep{garnett90,stasinska90,stasinska96,stasinska03}, is to compute each temperature from: i.e., each temperature is an ion and weighted volume mean."
+ llowever. the above three temperatures can also be obtained from their respective line [τιν ratios. ην II]. R[O H]. or ΠΙΟ HI]. as explained below.," However, the above three temperatures can also be obtained from their respective line flux ratios, $R$ [N II], $R$ [O II], or $R$ [O III], as explained below."
+ This second method has the advantage of being consistent with the observational way of obtaining7).. i.e. [vom line flux ratios.," This second method has the advantage of being consistent with the observational way of obtaining, i.e., from line flux ratios."
+ Ideally. temperature parameterizations should be independent of the method employed to compute (he individual temperatures.," Ideally, temperature parameterizations should be independent of the method employed to compute the individual temperatures."
+ We will compare temperatures computed with different methocs further on. but first. we will explain how we determined ).. ).. and from (heir respective line [lux ratios.," We will compare temperatures computed with different methods further on, but first, we will explain how we determined , , and from their respective line flux ratios."
+"— LA) =sind | A)) | — | Aj) =-sin@), | À)) | Finally. the Cartesian p-coordinates of the incoming particles position and velocity are used. in the following geometric relations to find the components in spherical Coordinates: = l(zyfr) Pian: ( tir,——- ES E εαν y, with. ría=Cro2dus)2ALτ","= + ) = + ) + = + ) = - + ) + Finally, the Cartesian $\bmath{p}$ -coordinates of the incoming particle's position and velocity are used in the following geometric relations to find the components in spherical coordinates: r = = (z_p / r) = 2 ( ) = ( x_p + y_p + z_p ) = - ( r - z_p ) = ( x_p - y_p ) with $r_{p0} = (x_p^2 + y_p^2)^{1/2}$."
+"ν The critical impact parameter bog depends on tje particle speed ey=eig, and cosAy=píidzssdi (eq. 30)).", The critical impact parameter $b_{\mathrm{crit}}$ depends on the particle speed $v_1 = v_{\mathrm{th}} u_1$ and $\cos \theta_0 = - v_1^{-1} \mathrm{d}z_p/\mathrm{d}t$ (eq. \ref{eq:b_crit}) ).
+ WbXbaa. then we integrate the equations of motion (13)) (15)) to determine where on the erain surface the particle hits.," If $b \le b_{\mathrm{crit}}$, then we integrate the equations of motion \ref{eq:r_dot_dot}) \ref{eq:phi_dot_dot}) ) to determine where on the grain surface the particle hits."
+ We have triecb various values of rig., We have tried various values of $r_{\mathrm{big}}$.
+ OF course. larger values vield higher accuracy but also require smaller tinx steps.," Of course, larger values yield higher accuracy but also require smaller time steps."
+ We found tiab rp.=5060 vields high accuracy ancl is not prohibitively: ime consuming., We found that $r_{\mathrm{big}} = 50 a$ yields high accuracy and is not prohibitively time consuming.
+ Substituting the collisional charging procedure described here in our charging simulations (and including both electrons. and protons). we examined a perlectls insulating erain with e0.1pm.," Substituting the collisional charging procedure described here in our charging simulations (and including both electrons and protons), we examined a perfectly insulating grain with $a = 0.1 \ \micron$."
+" With a duration of 100vr. we found that |p.layfea22.5 and. mi, ranges from 5.8.10 4106.010 Las cos6,4, ranges [rom 0 to 1."," With a duration of $100 \ \mathrm{yr}$, we found that $|p_z|_{\mathrm{av}}/ea \approx 2.5$ and $\tau_{\mathrm{flip}}$ ranges from $5.8 \times 10^{-4}$ to $6.0 \times 10^{-4}$ as $\cos \theta_{Jv}$ ranges from 0 to 1."
+ Phe are very close othe results obtained previously. neglecting erain drift (and ignoring protons) in the treatment of collisional charging (see ‘Table 1).," These are very close tothe results obtained previously, neglecting grain drift (and ignoring protons) in the treatment of collisional charging (see Table \ref{tab:sim}) )."
+" Due to precession of J about DB. 1e angle 6), changes on a time-scale short compared: witi the simulation time of 100vr (but an order of magnitude longer than ziii)."," Due to precession of $\bmath{J}$ about $\bmath{B}$, the angle $\theta_{Jv}$ changes on a time-scale short compared with the simulation time of $100 \ \mathrm{yr}$ (but an order of magnitude longer than $\tau_{\mathrm{flip}}$ )."
+" However. our results imply that the behavior of the electric dipole moment. is insensitive to he value of 8;,.."," However, our results imply that the behavior of the electric dipole moment is insensitive to the value of $\theta_{Jv}$ ."
+ Phus. we conclude that the neglect of erain drift in 85.does not vield significant error.," Thus, we conclude that the neglect of grain drift in \ref{sec:results} does not yield significant error."
+of the ULIRGs with multiple nuclei is ~20%.,of the ULIRGs with multiple nuclei is $\sim 20\%$.
+" In addition, Cuietal.(2001) have analyzed the nine samples of Suraceetal. in J-band."," In addition, \cite{cui01} have analyzed the nine samples of \cite{sur98} in $I$ -band."
+" Two ULIRGs are found to have multiple (1998)nuclei, i.e., the fraction is consistent with the results of Borneetal.(2000) and Cuietal. (2001)."," Two ULIRGs are found to have multiple nuclei, i.e., the fraction is consistent with the results of \cite{bor00} and \cite{cui01}."
+". On the other hand, Bushouseetal.(2002) and Veilleuxal. have claimed that the fraction of ULIRGs with multiple(2002) nuclei is less than 5%, based on their analysis of the data observed in near IR bands."," On the other hand, \cite{bus02} and \cite{vei02} have claimed that the fraction of ULIRGs with multiple nuclei is less than $5\%$, based on their analysis of the data observed in near IR bands."
+" Using theHST H-band data, Bushouseetal.(2002) reanalyzed 27 samples randomly selected from the 123 J-band samples, which include a part of Borne’s and Cui’s samples with multiple nuclei."," Using the$H$ -band data, \cite{bus02} reanalyzed 27 samples randomly selected from the 123 $I$ -band samples, which include a part of Borne's and Cui's samples with multiple nuclei."
+" In their analysis, only one sample is classified as a ULIRG with multiple nuclei."," In their analysis, only one sample is classified as a ULIRG with multiple nuclei."
+ The other sample includes some ULIRGs that have been already classified as ULIRGs with multiple nuclei by J-band image analysis (Borneetal.2000;Cui," The other sample includes some ULIRGs that have been already classified as ULIRGs with multiple nuclei by $I$ -band image analysis \citep{bor00,cui01}."
+ The reasons for excluding them from ULIRGs with 2001)..multiple nuclei are as follows: (1) the nucleus seems to locate on the tidal arm and is thought to be the tidal arm or multiple nuclei morphology do not appear in the H-band(2) images., The reasons for excluding them from ULIRGs with multiple nuclei are as follows: (1) the nucleus seems to locate on the tidal arm and is thought to be the tidal arm or (2) multiple nuclei morphology do not appear in the $H$ -band images.
+" Furthermore, Veilleuxetal.(2002) analyzed the K'""-band data of 118 ULIRGs taken by the Keck observatory."," Furthermore, \cite{vei02} analyzed the $K'$ -band data of 118 ULIRGs taken by the Keck observatory."
+ Their samples also include candidates of ULIRGs with multiple nuclei in HST J-band data., Their samples also include candidates of ULIRGs with multiple nuclei in HST $I$ -band data.
+ Some of them are excluded since multiple cores are not apparent in K’-band in spite of their appearance in I-, Some of them are excluded since multiple cores are not apparent in $K'$ -band in spite of their appearance in $I$ -band.
+" As a result, only 5 ULIRGs (~ 4%) are classified as ULIRGs with multiple nuclei."," As a result, only 5 ULIRGs $\sim 4\%$ ) are classified as ULIRGs with multiple nuclei."
+" In order to understand the difference between the results obtained using images taken in different bands, we made I-, H-, and K-band images of our simulation data usingPÉGASE."," In order to understand the difference between the results obtained using images taken in different bands, we made $I$ -, $H$ -, and $K$ -band images of our simulation data using."
+" We assumed that galaxies were at z~0.1 (1""2 kpc).", We assumed that galaxies were at $z \sim 0.1$ $1'' \sim 2~{\rm kpc}$ ).
+" The point spread function was assumed to be Gaussian, and its dispersion was calculated by o=Az/v21n2, where Ax is the spatial resolution."," The point spread function was assumed to be Gaussian, and its dispersion was calculated by $\sigma = \Delta x / \sqrt{2 \ln 2}$, where $\Delta x$ is the spatial resolution."
+ The spatial resolutions of J- and H-band images were given by the diffraction limit since we assumed the observations byHST., The spatial resolutions of $I$ - and $H$ -band images were given by the diffraction limit since we assumed the observations by.
+" The spatial resolutions of J- and H-bands are 0"".1 and 0"".2, respectively."," The spatial resolutions of $I$ - and $H$ -bands are $0''.1$ and $0''.2$, respectively."
+" For K-band images, the seeing determines the spatial resolution since we assumed ground-based observations by the University of Hawaii 2.2m telescope or the Keck spectroscopy."," For $K$ -band images, the seeing determines the spatial resolution since we assumed ground-based observations by the University of Hawaii $2.2~{\rm m}$ telescope or the Keck spectroscopy."
+" The spatial resolution, therefore, is set to be 0"".5 (Kimetal.2002)."," The spatial resolution, therefore, is set to be $0''.5$ \citep{kim02}."
+". For comparison, we also made the emulated images with K-band byTelescope (JWST) or ground-based 8m telescopes with adaptive optics (AO)."," For comparison, we also made the emulated images with $K$ -band by ) or ground-based $8$ m telescopes with adaptive optics (AO)."
+" We assume that the angular resolution of these future instruments is ~0"".05.", We assume that the angular resolution of these future instruments is $\sim 0''.05$.
+" Figure 8 shows I-, H-, and K-band images, in which dust extinction are not taken into account, of runs Hpp and HT."," Figure \ref{multibands} shows $I$ -, $H$ -, and $K$ -band images, in which dust extinction are not taken into account, of runs $H_{\rm{PP}}$ and $H_{\rm{TT}}$."
+" In I-band images, multiple core-like structures are clearly visible, while such cores are blurred in other bands due to the limited resolution and the contribution of the luminosity of old stars."," In $I$ -band images, multiple core-like structures are clearly visible, while such cores are blurred in other bands due to the limited resolution and the contribution of the luminosity of old stars."
+" In H-band images, cores are connected one another and form arm-like structures."," In $H$ -band images, cores are connected one another and form arm-like structures."
+" As a result, the cores seem to be in arms."," As a result, the cores seem to be in arms."
+" In K- images, multiple core structures are not resolved at all."," In $K$ -band images, multiple core structures are not resolved at all."
+" As a result, ULIRGs seem to have only single or double cores."," As a result, ULIRGs seem to have only single or double cores."
+" Thus, multiple cores were identified as a single or double cores in H- and K'-bands."," Thus, multiple cores were identified as a single or double cores in $H$ - and $K'$ -bands."
+ Our analysis indicates that ULIRGs with multiple cores can be misclassified as ones with single or double cores., Our analysis indicates that ULIRGs with multiple cores can be misclassified as ones with single or double cores.
+ The apparent discrepancy in the observational fractions of ULIRGs with multiple nuclei might have been caused by this difference in the spatial resolution., The apparent discrepancy in the observational fractions of ULIRGs with multiple nuclei might have been caused by this difference in the spatial resolution.
+" In the rightmost column in Fig. 8, "," In the rightmost column in Fig. \ref{multibands}, ,"
+we show the emulated K-band images of our simulations with the angular resolution of 0”.05., we show the emulated $K$ -band images of our simulations with the angular resolution of $0''.05$.
+" We can clearly see the multiple brightsources in these K-band images that were not visible in grand-base observations, since the resolution of the image is remarkably improved."," We can clearly see the multiple brightsources in these $K$ -band images that were not visible in grand-base observations, since the resolution of the image is remarkably improved."
+" We expect that, in the near future, observations of (U)LIRGs byJWST or ground-based 8m telescopes with AO will settle the problem of the band-to-band difference in the fraction of ULIRGs with multiple nuclei."," We expect that, in the near future, observations of (U)LIRGs by or ground-based $8$ m telescopes with AO will settle the problem of the band-to-band difference in the fraction of ULIRGs with multiple nuclei."
+" Figure 9 shows the observed I-, H-, and K-band images taking dust extinction into account."," Figure \ref{multibands_dust} shows the observed $I$ -, $H$ -, and $K$ -band images taking dust extinction into account."
+" Here, the cross-sections per H atom at H- and K-bands are assumed to be σηη=1.0x10733 and CHK=7.0x 10-2, respectively (Draine2003)."," Here, the cross-sections per H atom at $H$ - and $K$ -bands are assumed to be $\sigma _{H,\rm{H}}=1.0\times 10^{-22}$ and $\sigma _{H,\rm{K}}=7.0\times 10^{-23}$ , respectively \citep{dra03}."
+. We adopt qq=0.5 as a fiducial value., We adopt $q_d = 0.5$ as a fiducial value.
+" In the J-band image, the hypermassive star clusters and the galactic cores are obscured because of dust extinction although the resolution is enough to observe them."," In the $I$ -band image, the hypermassive star clusters and the galactic cores are obscured because of dust extinction although the resolution is enough to observe them."
+" On the other hand, in the H- and ground- K-band images, dust extinction becomes weaker although the clusters and the galactic cores are not resolved."," On the other hand, in the $H$ - and ground-based $K$ -band images, dust extinction becomes weaker although the clusters and the galactic cores are not resolved."
+" In the K-band image by or ground-based 8m telescopes with AO, the clusters and the galactic cores are clearly observable since not only dust extinction is small but also they are resolved."," In the $K$ -band image by or ground-based $8$ m telescopes with AO, the clusters and the galactic cores are clearly observable since not only dust extinction is small but also they are resolved."
+" When dust extinction is significant in J-band, we need the observationsby JWST or ground-based 8m telescopes with AO."," When dust extinction is significant in $I$ -band, we need the observationsby or ground-based $8$ m telescopes with AO."
+ We have performed high resolution N-body/SPH simulations of merging galaxies., We have performed high resolution $N$ -body/SPH simulations of merging galaxies.
+ We found that hypermassive star clusters with ~10°Mo form from disturbed gas disks in the central region (~ kpc).In, We found that hypermassive star clusters with $\sim 10^{8}~M_{\odot}$ form from disturbed gas disks in the central region $\sim \rm{kpc}$ ).In
+The Ηόόα line to continuum ratio reported by n-Pintado et ((1993) is typically about20%.,The $\alpha$ line to continuum ratio reported by n-Pintado et (1993) is typically about.
+". In contrast, the line peak temperature found here (25 mJy) is only about of the underlying continuum."," In contrast, the line peak temperature found here (25 mJy) is only about of the underlying continuum."
+ A somewhat more robust line to continuum ratio of about is obtained by averaging over the central 90 km | of the line., A somewhat more robust line to continuum ratio of about is obtained by averaging over the central 90 km $^{-1}$ of the line.
+" In LTE conditions and adopting a full width at half power of 90 kms ! and an electron temperature of 8.000 K, the H66« line is predicted to represent about of the underlying continuum (Altenhotf et 1981: euez 1982; n-Pintado 2002) in excellent agreement with the observed values."," In LTE conditions and adopting a full width at half power of 90 km $^{-1}$ and an electron temperature of 8,000 K, the $\alpha$ line is predicted to represent about of the underlying continuum (Altenhoff et 1981; guez 1982; n-Pintado 2002) in excellent agreement with the observed values."
+ We conclude that the H66a linc in MWC 349A is largely thermally excited and is not strongly masing as had been suggested by n-Pintado et (1993)., We conclude that the $\alpha$ line in MWC 349A is largely thermally excited and is not strongly masing as had been suggested by n-Pintado et (1993).
+" One could argue, however, that our results and those of n-Pintado et ((1993) are not necessarily incompatible because the two sets of observations might trace different regions of MWC 349A. In particular, the present observations were obtained in the most compact (D) configuration of the EVLA and are, therefore, sensitive to even fairly extended emission."," One could argue, however, that our results and those of n-Pintado et (1993) are not necessarily incompatible because the two sets of observations might trace different regions of MWC 349A. In particular, the present observations were obtained in the most compact (D) configuration of the EVLA and are, therefore, sensitive to even fairly extended emission."
+ If it existed. such extended emission would have been filtered out in the," If it existed, such extended emission would have been filtered out in the"
+and suggestions.,and suggestions.
+ MJV. also thanks the Brazilian agency FAPESP (Proc., MJV also thanks the Brazilian agency FAPESP (Proc.
+ No., No.
+ 96/00677-3) and PROPP/UESC for finantial support., 96/00677-3) and PROPP/UESC for finantial support.
+ VJP and RO thank the federal Brazilian agency CNPq for partial support., VJP and RO thank the federal Brazilian agency CNPq for partial support.
+ The authors would also like to thank the project PRONEX (41.96.0908.00) [or partial support., The authors would also like to thank the project PRONEX (41.96.0908.00) for partial support.
+versus the sound mode.,versus the sound mode.
+ To clarify. we emphasize that our 7 is the same quantity that is computed in |12].. but is not the same as the τι of Mülller. Israel and Stewart.," To clarify, we emphasize that our $\tau$ is the same quantity that is computed in \cite{conformalrelax1}, but is not the same as the $\tau_\Pi$ of Mülller, Israel and Stewart."
+ Llowever. il is necessary input lor the higher order relativistic dissipative fIuid dynamic description of the flow of matter and radiation.," However, it is necessary input for the higher order relativistic dissipative fluid dynamic description of the flow of matter and radiation."
+ We shall now derive a formula lor 7., We shall now derive a formula for $\tau$.
+" The lowest order solution. which was all (hat was needed to determine the shear viscosity. was (p) = ay dr!πα... BY) = —iDa, xod!(4"," The lowest order solution, which was all that was needed to determine the shear viscosity, was (r) = a_1 _r dr', (r) = -iD a_1 _r dr'."
+0) substituting (hese into Eqs. (3))-(3)), Substituting these into Eqs. \ref{hyzA}) \ref{hyzB}) )
+" leads to the next higher order terms. A) Sarg dpt ""E di (o) LY Bop) = by dr b, —iD dil gh) Amos(41)Dirichlet boundary conditions at infinity require da=bo0."," leads to the next higher order terms, (r) = a_2 + dr' a_3 + dr” (r”) (r”), (r) = b_2 + dr' b_3 - iD dr” (r”) (r”).Dirichlet boundary conditions at infinity require $a_2 = b_2 = 0$."
+ Finally. substitution of these solutions into Eq. (3))," Finally, substitution of these solutions into Eq. \ref{hyzC}) )"
+ leads to Der = — De +i(42) To complete thecaleulation the boundary. condition on the stretched horizon (3)) must be satisfied., leads to D^2 = - D + i. To complete thecalculation the boundary condition on the stretched horizon \ref{stretchedBC}) ) must be satisfied.
+ To this order it is, To this order it is (r-r_0) = - ( - i D )
+subject headings: galaxies;spiral — galaxies:IM galaxies:evolution ealaxies:siruciure — interstellarinolecules — interstellar:kinematics and dynamics Bars exert gravitational torques on the gas in the disks of spiral galaxies. resulting in eas inflow towards (he center (Quillen et al.,"Subject headings: galaxies:spiral — galaxies:ISM — galaxies:evolution — galaxies:structure — interstellar:molecules — interstellar:kinematics and dynamics Bars exert gravitational torques on the gas in the disks of spiral galaxies, resulting in gas inflow towards the center (Quillen et al."
+ 1995: Reean. Vogel. Teuben 1997: ltegan. Sheth. Vogel 1999).," 1995; Regan, Vogel, Teuben 1997; Regan, Sheth, Vogel 1999)."
+ This results in a significant increase in the gas mass in the center ol a bar (Sakamoto et al., This results in a significant increase in the gas mass in the center of a bar (Sakamoto et al.
+ 1999: Sheth 2001). often leading to increased star formation and even starburst activity in the nucleus (IIo. Filippenko. Sargent. 1997: Jogee. Ixennev.," 1999; Sheth 2001), often leading to increased star formation and even starburst activity in the nucleus (Ho, Filippenko, Sargent, 1997; Jogee, Kenney,"
+ , 
+resulting mean LL column density is larger for a fractal mecium. where the highest density regions are more cillicult to ionise and recombine faster.,"resulting mean HI column density is larger for a fractal medium, where the highest density regions are more difficult to ionise and recombine faster."
+" Also with an increased stellar surface density (runs € and D) the mean Ny; in a EDD remains as high as zz2.7107"" 7. while in a GDD it is z5.9«1070 em27. with. values as low as 2103T 7. Do"," Also with an increased stellar surface density (runs C and D) the mean $N_{HI}$ in a FDD remains as high as $\approx 2.7 \times 10^{20}$ $^{-2}$, while in a GDD it is $\approx 5.9 
+\times 10^{19}$ $^{-2}$, with values as low as $\approx 10^{17}$ $^{-2}$."
+cThe reduced ionisation fraction in the case of an inhomogeneous mecium was also noted in the mocel for the cliffusect ionised, The reduced ionisation fraction in the case of an inhomogeneous medium was also noted in the model for the diffused ionised
+"The discovery Chat the hot Jupiter ID 209453b is losing mass was rather unexpected (Recently, Lecavelier des Etangs et al. (","The discovery that the hot Jupiter HD 209458b is losing mass was rather unexpected (Recently, Lecavelier des Etangs et al. ("
+2010) have found atmospheric evaporation in IID 1891733b.,2010) have found atmospheric evaporation in HD 189733b.
+ It is the second extrasolar planet whose mass loss has been detected.)., It is the second extrasolar planet whose mass loss has been detected.).
+ The excess absorption in Lyman-alpha first found by Vidal-Madjar et al. (, The excess absorption in Lyman-alpha first found by Vidal-Madjar et al. (
+2003) (VM02) and later confirmed by Linsky et al. (,2003) (VM03) and later confirmed by Linsky et al. (
+2010) could be explained either bv mass loss of the atmosphere due to XUV energy input [rom host stars (Lamuner οἱ al.,2010) could be explained either by mass loss of the atmosphere due to XUV energy input from host stars (Lammer et al.
+ 2003: Lecavelier des Etanes et al., 2003; Lecavelier des Etangs et al.
+ 2004: Yelle 2004. 2006: Tian et al.," 2004; Yelle 2004, 2006; Tian et al."
+ 2005: Garcia. Munoz 2007: Penz οἱ al., 2005; Garcia Munoz 2007; Penz et al.
+ 2008: Miurav-clay et al., 2008; Murray-clay et al.
+ 2009: Lammer et al., 2009; Lammer et al.
+ 2009) or bv charge exchange between the stellar wind and the planetary. escaping exosphere (lolimstronun et al., 2009) or by charge exchange between the stellar wind and the planetary escaping exosphere (Holmströmm et al.
+ 2008)., 2008).
+ For the former. all of these models describe the thermal particle escape.," For the former, all of these models describe the thermal particle escape."
+ For the latter. Exkaev et al. (," For the latter, Erkaev et al. ("
+2005) and Holmstrom et al. (,2005) and Holmströmm et al. (
+2008) brought forward loss of nonthermal neutral atoms due to interaction between the stellar wind aud the exosphere (more details see Ekenbàcek οἱ al.,2008) brought forward loss of nonthermal neutral atoms due to interaction between the stellar wind and the exosphere (more details see Ekenbäcck et al.
+ 2010)., 2010).
+ The model of Erkaev et al. (, The model of Erkaev et al. (
+2005) underestimates evidently the particle loss rates. but. Ekenbacck et al. (,"2005) underestimates evidently the particle loss rates, but Ekenbäcck et al. ("
+2010) modeled the production of neutral hydrogen and match the feature of £ja absorption.,2010) modeled the production of neutral hydrogen and match the feature of $Ly\alpha$ absorption.
+ It is not easy to distinguish which process dominates Lya absorption more., It is not easy to distinguish which process dominates $Ly\alpha$ absorption more.
+ Den-Jaffel ILIosseini (2010) found that either energetic III of stellar origin or thermal III populations in the planetary atmosphere could fit Lya observations., Ben-Jaffel Hosseini (2010) found that either energetic HI of stellar origin or thermal HI populations in the planetary atmosphere could fit $Ly\alpha$ observations.
+ Ixoskinen et al. (, Koskinen et al. (
+2010)(IX10) used an empirical model to analyze UV transit depths. and their results showed that observations can be explained solely by absorption in the upper alinosphere while the process ol charge exchange could not be necessary.,"2010)(K10) used an empirical model to analyze UV transit depths, and their results showed that observations can be explained solely by absorption in the upper atmosphere while the process of charge exchange could not be necessary."
+ It is also noteworthy (hat the properties of planetary magnetic field are unclear until now., It is also noteworthy that the properties of planetary magnetic field are unclear until now.
+ Thus. it is important for both thermal models and nonthermal models to calculate sell-consistentlv the deflection distance around," Thus, it is important for both thermal models and nonthermal models to calculate self-consistently the deflection distance around"
+position of V802 Tau NE.,position of V892 Tau NE.
+ The absolute measurement accuracy of the ppoiutiug is αρα UUsers Handbook). so the offset of the source iu the EPIC-pu nage from the radio coordinates of V892 Tau is well within 1 6 of the pointing accuracy.," The absolute measurement accuracy of the pointing is 4 arcsec Users' Handbook), so the offset of the source in the EPIC-pn image from the radio coordinates of V892 Tau is well within 1 $\sigma$ of the pointing accuracy."
+ This is fully consistent with V892 Tau being the dominant source of N-ray cluission throughout the oobservation. as it is the case during the oobservation. where V892 Tau is roughly 10 times more huninous than V892 Tau NE.," This is fully consistent with V892 Tau being the dominant source of X-ray emission throughout the observation, as it is the case during the observation, where V892 Tau is roughly 10 times more luminous than V892 Tau NE."
+ Second. the pposition of the source associated with the Svsteni V892 Tau before and durius the flare is constant within 0.7aresoet.. while the angular separation of V802 Tan and V892 Tau NE is L1 aresec.," Second, the position of the source associated with the system V892 Tau before and during the flare is constant within 0.7, while the angular separation of V892 Tau and V892 Tau NE is 4.1 arcsec."
+ At the same time (as shown in the previous section) the spectral characteristic of the EPIC-pu data on the V892 Tau system while quiescent are in good agreenmieut with the characteristics derived fron the spectral data on V892 Tau frou but sienificautly differcut from the oues derived from the sspectra data on V892 Tau NE., At the same time (as shown in the previous section) the spectral characteristic of the EPIC-pn data on the V892 Tau system while quiescent are in good agreement with the characteristics derived from the spectral data on V892 Tau from but significantly different from the ones derived from the spectral data on V892 Tau NE.
+ These two facts stronely iucdicate hat the ddata are. for both the quiescent emission aud the flare. fully dominated by he cussion of V892 Tau.," These two facts strongly indicate that the data are, for both the quiescent emission and the flare, fully dominated by the emission of V892 Tau."
+ Πας the flare been associate with V892 Tau NE. a shift of the N-ray source position durus the Hare would have been observed.," Had the flare been associated with V892 Tau NE, a shift of the X-ray source position during the flare would have been observed."
+ Finally. if one wanted to explain the ]lheht curves as mostly determined by V892 Tau NE. disregarding the evidence coming roni the pointing aud spectral information. one would have to require that the source which does not appear significautly variable in the observation becomes from ten (before the fiue) to a hundred times (dunug the fare) brighter during oobservation.," Finally, if one wanted to explain the light curves as mostly determined by V892 Tau NE, disregarding the evidence coming from the pointing and spectral information, one would have to require that the source which does not appear significantly variable in the observation becomes from ten (before the flare) to a hundred times (during the flare) brighter during observation."
+ This is uulikelsy., This is unlikely.
+ Given the facts above. we conclude that the ]lheht curve shown iu Fie.," Given the facts above, we conclude that the light curve shown in Fig."
+ 1 and the EPIC-pu spectra shown in Fig., \ref{fig:lc_xmm} and the EPIC-pn spectra shown in Fig.
+ 6 aud Fig., \ref{fig:ps_xmm} and Fig.
+ 5 are determined by the activity of the Herbig Ac star V892 Tan and minimally iuflueuced by the presence of V892 Tau NE., \ref{fig:ps_xmm74ks} are determined by the activity of the Herbig Ae star V892 Tau and minimally influenced by the presence of V892 Tau NE.
+ Cousequeutly. below we discuss the ddata as originating frou Vso2 Tau aud do not further cohunent on the uuiresolve presence of V892 Tau NE.," Consequently, below we discuss the data as originating from V892 Tau and do not further comment on the unresolved presence of V892 Tau NE."
+Cataclysmic variables (CVs) are thought to form. from detached white dwarf (WD) / main-sequence star (MS) binaries that have experienced a common-envelope (CE) phase 2008).,Cataclysmic variables (CVs) are thought to form from detached white dwarf (WD) / main-sequence star (MS) binaries that have experienced a common-envelope (CE) phase .
+" In these ""post-common-envelope binaries” (PCEBs). the white dwarf represents the more massive component in the system and is therefore called the primary. while the usually late-type (K—M) main-sequence star is known as the secondary."," In these “post-common-envelope binaries"" (PCEBs), the white dwarf represents the more massive component in the system and is therefore called the primary, while the usually late-type (K--M) main-sequence star is known as the secondary."
+ After the end of the CE phase. angular-momentum loss due to magnetic braking and/or gravitational radiation continues to decrease the separation between the two stars.," After the end of the CE phase, angular-momentum loss due to magnetic braking and/or gravitational radiation continues to decrease the separation between the two stars."
+ This eventually brings the secondary's Roche lobe into contact with the stellar surface. thus initiating stable mass-transfer via Roche-lobe overflow and the semi-detached CV phasetherein).," This eventually brings the secondary's Roche lobe into contact with the stellar surface, thus initiating stable mass-transfer via Roche-lobe overflow and the semi-detached CV phase."
+. There is growing evidence that the aceretion of material from the secondary star does not start. with Roche-lobe overflow., There is growing evidence that the accretion of material from the secondary star does not start with Roche-lobe overflow.
+ The discovery of so-called “low accretion rate polars”2002).. which are likely progenitors of magnetic CVs2007). and the detection of metal absorption lines in the UV spectra of non-magnetic PCEBs2008).. indicate that wind accretion from the active chromosphere of the secondary star is à common phenomenon in detached. short-period. WD/MS binaries.," The discovery of so-called “low accretion rate polars”, which are likely progenitors of magnetic CVs, and the detection of metal absorption lines in the UV spectra of non-magnetic PCEBs, indicate that wind accretion from the active chromosphere of the secondary star is a common phenomenon in detached, short-period, WD/MS binaries."
+ found the Ha emission line in the detached PCEB to be a combinatior of two components., found the $\alpha$ emission line in the detached PCEB to be a combination of two components.
+ The radial velocity variations in the stronger He component agree well with those shown by the TIO absorption. and was thus identified as originating in the secondary star.," The radial velocity variations in the stronger $\alpha$ component agree well with those shown by the TiO absorption, and was thus identified as originating in the secondary star."
+ The radial velocities of the weaker component showed an — within the errors — anti-phased behaviour with respect to the secondary star exhibiting a lower amplitude. and therefore had to be produced on the side of the centre-of-mass opposite to the secondary.," The radial velocities of the weaker component showed an – within the errors – anti-phased behaviour with respect to the secondary star exhibiting a lower amplitude, and therefore had to be produced on the side of the centre-of-mass opposite to the secondary."
+ However. the low spectral resolution. of the data. and the contamination of the white-dwarf Balmer absorption lines with emission cores from the secondary. impeded velocity measurements of other. unambiguously intrinsic. spectral features from the white dwarf.," However, the low spectral resolution of the data, and the contamination of the white-dwarf Balmer absorption lines with emission cores from the secondary, impeded velocity measurements of other, unambiguously intrinsic, spectral features from the white dwarf."
+ Thus. the true origin of the weak emission component remained unresolved.," Thus, the true origin of the weak emission component remained unresolved."
+ The low temperature of the white dwarf in LTT 560 -7500K. Paper D and the evidence of ongoing accretion imply that there is a high probability that there are narrow metal absorption lines in the white-dwarf spectrum2003).," The low temperature of the white dwarf in LTT 560 $T \sim 7500~\mathrm{K}$, Paper I) and the evidence of ongoing accretion imply that there is a high probability that there are narrow metal absorption lines in the white-dwarf spectrum."
+. The radial velocity curve of such lines should be undisturbed and thus faithfully track the motion of the white dwarf. motivating the high-resolution study presented in this paper.," The radial velocity curve of such lines should be undisturbed and thus faithfully track the motion of the white dwarf, motivating the high-resolution study presented in this paper."
+ LTT 560 was observed on August 16. 2007 using UVES on UTI (ANTU) at ESO Paranal.," LTT 560 was observed on August 16, 2007 using UVES on UT1 (ANTU) at ESO Paranal."
+" A total of 38 Echelle spectra was taken over 7.3 h in one blue (3260-4528 A) and two red (5681- A. 7661-9439A. hereafter ""lower"" and “upper” red spectrum. respectively) spectral ranges. with resolving power - 700000."," A total of 38 Echelle spectra was taken over 7.3 h in one blue (3260–4528 ) and two red (5681-7519 , 7661-9439, hereafter ""lower"" and ""upper"" red spectrum, respectively) spectral ranges, with resolving power $\sim$ 000."
+ The exposure time per spectrum was 600 s. which corresponds to roughly 0.05 orbital cycles.," The exposure time per spectrum was 600 s, which corresponds to roughly 0.05 orbital cycles."
+ The sequence of time-resolved spectra had to be interrupted for about 45 min when the object was close to the zenith. where the VLT cannot observe.," The sequence of time-resolved spectra had to be interrupted for about 45 min when the object was close to the zenith, where the VLT cannot observe."
+ The data reduction was performed using Gasgano and the UVES pipeline (version 3.9.0) in a step-by-step mode., The data reduction was performed using Gasgano and the UVES pipeline (version 3.9.0) in a step-by-step mode.
+" This included bias and flat correction, wavelength calibration with a ThAr lamp. and flux calibration using the standard. star LTT 7987."," This included bias and flat correction, wavelength calibration with a ThAr lamp, and flux calibration using the standard star LTT 7987."
+ The response function proved rather unsatisfactory. with the flux-calibrated spectra still containing several “order wriggles”.," The response function proved rather unsatisfactory, with the flux-calibrated spectra still containing several ""order wriggles""."
+ However. since this bears no relevance to the work described in this paper. we did not attempt to resolve this issue.," However, since this bears no relevance to the work described in this paper, we did not attempt to resolve this issue."
+from the 2\LASS survey.,from the 2MASS survey.
+ Finally. for 52 stars. Strómmegren photometry is available from: NorelstrGnun ct al (2004) and Tyeho photometry from ESA (1997).," Finally, for 52 stars, Strömmgren photometry is available from Nordströmm et al (2004) and Tycho photometry from ESA (1997)."
+" For cach colour. the clleetive temperatures of the stars are fit to a linear (first. order). function. of colour and ‘Ve1I] (we also tried second. order fits. but these were not significant improvements on linear fitting): in practice. it was convenient to fit Gop,=5040/1.yy rather than Zipp directly."," For each colour, the effective temperatures of the stars are fit to a linear (first order) function of colour and [Fe/H] (we also tried second order fits, but these were not significant improvements on linear fitting); in practice, it was convenient to fit $\theta_{\mathrm eff} = 5040/\Teff$ rather than $\Teff$ directly."
+ For example. for BoV. we fitted for (04.602.055) in the relation (BV) bas) .," For example, for $B-V$, we fitted for $(a_1,a_2,a_3)$ in the relation = a_0 + a_1 (B-V) + a_2 ."
+- The reason we fit for θες as a function of colour and metallicity. is because the relative errors in colour and Fe/11] are always much smaller than the relative error in θε.," The reason we fit for $\theta_{\mathrm eff}$ as a function of colour and metallicity, is because the relative errors in colour and [Fe/H] are always much smaller than the relative error in $\theta_{\mathrm eff}$."
+ Figure 2. shows the result of the fit., Figure \ref{tefit} shows the result of the fit.
+ The relations obtained can then be used to solve for the solar colours (e.g. (D.V). ). under the assumption of νετ=5777 d and Fe/l].= 0.0.," The relations obtained can then be used to solve for the solar colours (e.g. $(B-V)_\odot$ ), under the assumption of $\Teff,\odot = 5777$ K and $_\odot = 0.0$ ."
+ Formally. we have converted the relation to the form (c.g. for BV) Dbπμ by | by pp)) | be in order to make clearer the ellect of temperature and metallicity on derived colours.," Formally, we have converted the relation to the form (e.g. for $B-V$ ) B-V = b_0 + b_1 ) + b_2 in order to make clearer the effect of temperature and metallicity on derived colours."
+ Examination of column + in table 1.. for example. shows that the ellect of metallicity: in the colour transformations decreases toward veckeler colour bands. ancl reverses its sign in the Hi. as one would expect.," Examination of column 4 in table \ref{fits}, for example, shows that the effect of metallicity in the colour transformations decreases toward redder colour bands, and reverses its sign in the IR, as one would expect."
+ The parameters for each of the fitted colours (6j.be.bs) are shown in Table 1.. and the solar colours derived: are shown in Table 2..," The parameters for each of the fitted colours $(b_1,b_2,b_3)$ are shown in Table \ref{fits}, and the solar colours derived are shown in Table \ref{solcol}."
+ We have carefully propagated the errors in fitting in order to obtain 1-0 error estimates on the derived solar colours., We have carefully propagated the errors in fitting in order to obtain $\sigma$ error estimates on the derived solar colours.
+ The error quoted in Table 2. for each colour is the total dispersion of the fit., The error quoted in Table \ref{solcol} for each colour is the total dispersion of the fit.
+ In addition to the direct Tycho ByVpcolour shown in Table 2.. we also give the resulting Johnson colour derived: adopting the full transformation used for the D.V values given in ESA (1997).," In addition to the direct Tycho $B_{\rm T}-V_{\rm T}$colour shown in Table \ref{solcol}, we also give the resulting Johnson colour derived adopting the full transformation used for the $B-V$ values given in ESA (1997)."
+ The result is closely consistent. with the proper Johnson BV fit., The result is closely consistent with the proper Johnson $B-V$ fit.
+ We estimate the size of the systematic error as follows: the LIREM temperature scale is now known to be consistent with the stellar interferometry based. temperature scale to within 18 Ix. (cüspersion 62 Ix. standard error 20 WK: rez Molénndez 20052).," We estimate the size of the systematic error as follows: the IRFM temperature scale is now known to be consistent with the stellar interferometry based temperature scale to within 18 K (dispersion 62 K, standard error 20 K; rez Melénndez 2005a)."
+ This result is consistent. with zero olfset. but also with a systematic olfset of order 20 Ix. An olfset of order 20 Ix. produces systematic shifts in our derived. colours of less than half the internal error: had. we included this potential svstematic error (in quadrature). our error estimates in Table 2. would increase by or Less for FI colours.," This result is consistent with zero offset, but also with a systematic offset of order 20 K. An offset of order 20 K produces systematic shifts in our derived colours of less than half the internal error; had we included this potential systematic error (in quadrature), our error estimates in Table \ref{solcol} would increase by or less for all colours."
+ Interestingly. there are many more stars with accurate effective temperatures ancl metallicities for which xhotometry is available in the LR and Strómmgren colours. iin in CDVRE since the stars are so bright (and we »esentlv only have about a dozen stars in the sample) we are planning to extend the CDVI sample with observations a 1e 30 cm telescope on La Palma.," Interestingly, there are many more stars with accurate effective temperatures and metallicities for which photometry is available in the IR and Strömmgren colours, than in $UBVRI$ — since the stars are so bright (and we presently only have about a dozen stars in the sample) we are planning to extend the $UBVRI$ sample with observations at the 30 cm telescope on La Palma."
+" When we tried to extend our study to photometry in a similar wav as for the other colours. we found. tha rere is almost no overlap between Sloan Digital Sky Survey (SDSS) standard stars and stars with either HUEM Zep, or Vefl] measured."," When we tried to extend our study to photometry in a similar way as for the other colours, we found that there is almost no overlap between Sloan Digital Sky Survey (SDSS) standard stars and stars with either IRFM $\Teff$ or [Fe/H] measured."
+ Hopefully this unfortunate situation wil soon be remecdied. but in the meantime we apply a secondary method. to obtain the solar colours in the SDSS. bands.," Hopefully this unfortunate situation will soon be remedied, but in the meantime we apply a secondary method to obtain the solar colours in the SDSS bands."
+ We have usec the observationallv based: transformations between SDSS and JohnsonCousins photometry from Smith et al. (, We have used the observationally based transformations between SDSS and Johnson–Cousins photometry from Smith et al. (
+2002). Ixaraali. Dilir. Tunceel (2005) and Air. Waraali Tunceel (2005).,"2002), Karaali, Bilir, Tunçeel (2005) and Bilir, Karaali Tunçeel (2005)."
+ To give an estimate of the present uncertainty of the transformations. we have added athe dispersion between them to the propagated dispersion from the fit giving the total number shown in Table 2..," To give an estimate of the present uncertainty of the transformations, we have added the dispersion between them to the propagated dispersion from the fit giving the total number shown in Table \ref{solcol}."
+ We have derived colour estimates for the Sun from Sun-like stars. taking advantage of the fact that the adopted stellar ellective temperature scale from the HUEM has been recentIv confirmed as both accurate and. precise from interferometric measurements of stellar ciameters rez Moelénndez 2005a).," We have derived colour estimates for the Sun from Sun-like stars, taking advantage of the fact that the adopted stellar effective temperature scale from the IRFM has been recently confirmed as both accurate and precise from interferometric measurements of stellar diameters rez Melénndez 2005a)."
+ This is the primary reason for confidence in deriving solar colours from. Sun-like stars. but we have also restricted ourselves to the highest quality metallicity and colour cata as well.," This is the primary reason for confidence in deriving solar colours from Sun-like stars, but we have also restricted ourselves to the highest quality metallicity and colour data as well."
+ While undertaking thisstudy. we found that solar colours had been derived in a similar manner by irez Molénndez(2005b).," While undertaking thisstudy, we found that solar colours had been derived in a similar manner by rez Melénndez(2005b)."
+ Phey derived broad andmedium band colours : Johnson/Cousins. Vilnius. Strómmgren by. DDO. 2MASS and Tycho photometric systems. as part of a," They derived broad andmedium band colours : Johnson/Cousins, Vilnius, Strömmgren $b-y$ , DDO, 2MASS and Tycho photometric systems, as part of a"
+"If the fraction ε of the cloud mass is converted into stars after the cloud lifetime 7.1, the star formation rate, 7, is written as where Ma; is the total gas mass in the galaxy.","If the fraction $\epsilon$ of the cloud mass is converted into stars after the cloud lifetime $\tau_\mathrm{cl}$, the star formation rate, $\psi$, is written as where $M_\mathrm{gas}$ is the total gas mass in the galaxy."
+" Then, equation (29)) can be written by using thestar formation rate as where D=Maus:/Mgas is the dust-to-gas ratio."," Then, equation \ref{eq:dMdt2}) ) can be written by using thestar formation rate as where $\mathcal{D}\equiv M_\mathrm{dust}/M_\mathrm{gas}$ is the dust-to-gas ratio."
+ Here we summarize the recipe to include the grain growth by accretion into dust enrichment models., Here we summarize the recipe to include the grain growth by accretion into dust enrichment models.
+" There are two methods as described in sections 4.1 and 4.2,, called Method I and Method II, respectively."," There are two methods as described in sections \ref{subsec:dMdt1} and \ref{subsec:dMdt2}, called Method I and Method II, respectively."
+ We demonstrate that our results so far are really implemented into galaxy evolution models with dust enrichment., We demonstrate that our results so far are really implemented into galaxy evolution models with dust enrichment.
+ The aim here is not to construct an elaborate chemical evolution model but to focus on the effect of grain growth., The aim here is not to construct an elaborate chemical evolution model but to focus on the effect of grain growth.
+ We adopt a one-zone closed-box model of dust/metal enrichment by Hirashita(2000b) (seealsoDwek1998;Lisenfeld&Ferrara 1998).," We adopt a one-zone closed-box model of dust/metal enrichment by \citet{hirashita00b}
+ \citep[see also][]{dwek98,lisenfeld98}."
+". The model treats the evolutions of total gas, metals, and dust masses (Maas, Mz, and Maus. respectively) in the galaxy."," The model treats the evolutions of total gas, metals, and dust masses $M_\mathrm{gas}$, $M_{Z}$, and $M_\mathrm{dust}$, respectively) in the galaxy."
+" In this model, the metals include not only gas phase elements but also dust."," In this model, the metals include not only gas phase elements but also dust."
+" The equations are written as where E and Ez are the rate of the total injection of mass (gas + dust) and metal mass from stars, respectively, fin is the dust condensation efficiency of the metals in the stellar ejecta, and τον is the time-scale of dust destruction by SN shocks."," The equations are written as where $E$ and $E_Z$ are the rate of the total injection of mass (gas + dust) and metal mass from stars, respectively, $f_\mathrm{in}$ is the dust condensation efficiency of the metals in the stellar ejecta, and $\tau_\mathrm{SN}$ is the time-scale of dust destruction by SN shocks."
+" For simplicity, we do not treat individual element X but treat the entire metals; however, if we are interested in a specific dust-composing element X, we can replace the relevant quantities with ones specific for element X. The simple treatments above are sufficient to demonstrate that our scheme of grain growth in clouds for any grain size distribution is really applicable to the galaxy evolution models."," For simplicity, we do not treat individual element X but treat the entire metals; however, if we are interested in a specific dust-composing element X, we can replace the relevant quantities with ones specific for element X. The simple treatments above are sufficient to demonstrate that our scheme of grain growth in clouds for any grain size distribution is really applicable to the galaxy evolution models."
+ We refer to other papers (e.g.Dwek1998;Zhukovskaetal.2008;Galletal.2011a) for more detailed dust enrichment models.," We refer to other papers \citep[e.g.][]{dwek98,zhukovska08,gall11} for more detailed dust enrichment models."
+" Since we are not interested in the detailed history of dust production in stellar ejecta, we adopt the instantaneous recycling approximation; that is, a star with mm, (m is the age stellar mass, and m. is the turn-off mass at age t) dies instantaneously after its birth, leaving a remnant of mass w,,."," Since we are not interested in the detailed history of dust production in stellar ejecta, we adopt the instantaneous recycling approximation; that is, a star with $m>m_t$ $m$ is the zero-age stellar mass, and $m_t$ is the turn-off mass at age $t$ ) dies instantaneously after its birth, leaving a remnant of mass $w_m$."
+" Once the initial mass function (IMF) is fixed, the returned fraction of the mass fromformed stars, R, and the mass fraction of metals that is newly produced and ejected by stars, yz, are evaluated."," Once the initial mass function (IMF) is fixed, the returned fraction of the mass fromformed stars, $\mathcal{R}$ , and the mass fraction of metals that is newly produced and ejected by stars, $\mathcal{Y}_{Z}$, are evaluated."
+" Using these quantities, we write We adopt R=0.18 and Yx=0.013 (Appendix)."," Using these quantities, we write We adopt $\mathcal{R}=0.18$ and $\mathcal{Y}_\mathrm{X}=0.013$ (Appendix)."
+" For fin, we examine two cases: fin=0.1 and 0.01, which correspond to the fiducial and the lower efficiency cases in Inoue (2011),, respectively."," For $f_\mathrm{in}$ , we examine two cases: $f_\mathrm{in}=0.1$ and 0.01, which correspond to the fiducial and the lower efficiency cases in \citet{inoue11}, , respectively."
+" For the time-scale of dust destruction by SNe, we adopt an expression Τον= Meas/(€sMsy), where e,and M;are the dust destruction efficiency and the gas mass swept by a single"," For the time-scale of dust destruction by SNe, we adopt an expression $\tau_\mathrm{SN}=M_\mathrm{gas}/(\epsilon_\mathrm{s}
+M_\mathrm{s}\gamma )$ , where $\epsilon_\mathrm{s}$and $M_\mathrm{s}$are the dust destruction efficiency and the gas mass swept by a single"
+Migration of gas giant planets due to interactions wilh a circumstellar eas disk can play a major role in defining (he architecture of planetary. svstems.,Migration of gas giant planets due to interactions with a circumstellar gas disk can play a major role in defining the architecture of planetary systems.
+ Work on migration (seereviewbyPapaloizouetal.2007) has included gravitational interaction of planets with both laminar and turbulent disks., Work on migration \citep[see review by][]{papaloizou2007} has included gravitational interaction of planets with both laminar and turbulent disks.
+ Llowever. radiative transport. detailed equations of state (EOS). and the sell-gravitv of the gas disk have usually been ignored: the effects of a non-isolhermal EOS have only recently been included (e.g..Paardekooper&Mellema2006:Paardekooperοἱal.2010. 2011).," However, radiative transport, detailed equations of state (EOS), and the self-gravity of the gas disk have usually been ignored; the effects of a non-isothermal EOS have only recently been included \citep[e.g.,][]{paardekooper2006,paardekooper2010,paardekooper2011}."
+. Emerging studies. such as Baruteanetal.(2011) and the one reported here. are beeinning to address some of these issues.," Emerging studies, such as \citet{baruteau2011} and the one reported here, are beginning to address some of these issues."
+" Boss(2005). and Maveretal.(2004) examined radial migration of planet-mass fragments in gravitationallyv unstable disks. but their disks were violently disrupted bv fragmentation under conditions (radii < 40 AU. disk-to-star mass ratios Mj/AM,~0.1. and stellar mass A,~1AL.) where fragmentation may not actually occur (Ralikov2005.2007:Boleyetal.2006.20072:&Durisen2008:Forganetal.2009:Cai 2010)."," \citet{boss2005} and \citet{mayer2004} examined radial migration of planet-mass fragments in gravitationally unstable disks, but their disks were violently disrupted by fragmentation under conditions (radii $<$ 40 AU, disk-to-star mass ratios $M_d/M_s \sim 0.1$, and stellar mass $M_s \sim 1\;M_{\odot}$ ) where fragmentation may not actually occur \citep{rafikov2005,rafikov2007,boley2006,boley2007b,boley2008,forgan2009,cai2010}."
+. More recently. fragmentation into clumps wilh eas eiant or brown dwarl masses has been documented in nunerical simulations of disks Chat are relatively massive (AL)/\/. a few tenths) and spatially extended (outer radii > 50 AU) (Ixrumholzοἱal.2007:Stamatelloset&Whitworth2009:DolevBoleyοἱal. 2010).. where fragmentation is expected [rom senmi-analvtic arguments (e.g..Clarke2009:RafikovDoclson-Robinsonetal.2009).," More recently, fragmentation into clumps with gas giant or brown dwarf masses has been documented in numerical simulations of disks that are relatively massive $M_d/M_s \sim$ a few tenths) and spatially extended (outer radii $>$ 50 AU) \citep{krumholz2007,stamatellos2007,stamatellos2009,boley2009,boley2010}, where fragmentation is expected from semi-analytic arguments \citep[e.g.,][]{clarke2009,rafikov2009,dodson2009}."
+. The late of the clumps then depends in part on (their radial migration. which is a chaotic and messy affair in a fragmenting disk (e.g..Boley2009:Dolevetal.2010:Vorobvov&Basu2010a:&Durisen 2010)..," The fate of the clumps then depends in part on their radial migration, which is a chaotic and messy affair in a fragmenting disk \citep[e.g.,][]{boley2009,boley2010,vorobyov2010,boley2010b}."
+ The occurrence of gravitational instabilities (GIs) maa be episoclic (e.g..Basu2006.2010b:Zhuetal. 2010).," The occurrence of gravitational instabilities (GIs) may be episodic \citep[e.g.,][]{vorobyov2006,vorobyov2010b,zhu2010}."
+. Clumps that survive and contract to the dimensions of voung planets can later find (themselves in a disk that erupts again into GI activitv., Clumps that survive and contract to the dimensions of young planets can later find themselves in a disk that erupts again into GI activity.
+ As (he star/disk svstem evolves. such a protoplanet may end up in a region of a Gl-active disk where fragmentation does not occur.," As the star/disk system evolves, such a protoplanet may end up in a region of a GI-active disk where fragmentation does not occur."
+each patch.,each patch.
+ The covariance between the resulting patch correlation functions was measured as an estimator of the true covariance. in S angular separation bins logarithmically spaced from 1 to 90 and in 3 redshift bins (leading to 6 redshift pair bins).," The covariance between the resulting patch correlation functions was measured as an estimator of the true covariance, in 8 angular separation bins logarithmically spaced from $1'$ to $90'$ and in 3 redshift bins (leading to 6 redshift pair bins)."
+ The diagonal elements of the covariance matrix for mean correlation functions are measured to be approximately 10tt per square degree. making the sample covariance the dominant source of error for larger angles and higher redshifts for both our grouncd-based survey. with diagonal element values of 10.ο+? and Euclid. 10.LP.15.," The diagonal elements of the covariance matrix for mean correlation functions are measured to be approximately $10^{-11}-10^{-9}$ per square degree, making the sample covariance the dominant source of error for larger angles and higher redshifts for both our ground-based survey, with diagonal element values of $10^{-15}-10^{-13}$ and Euclid, $10^{-15}-10^{-13}$."
+ μονο should be compared with shape nolse covariance contributions of 101510.1. for eround-based and. 10l67 for Euclid., These should be compared with shape noise covariance contributions of $10^{-15}-10^{-11}$ for ground-based and $10^{-16}-10^{-12}$ for Euclid.
+" Phe covariances were included in our us estimations using (c£.Llartlapetal.2007) where d is the cdata here the Πάσα ACDA correlation function in redshift and angular separation bins: fois the alternative gravity. model correlation. function in those bins. m,=75 is the number of realisations of correlation functions usec in the calculation of Coo. ane my=48 is the total number of bins in angular separation an redshift."," The covariances were included in our $\chi^2$ estimations using \citep[c.f.][]{Hartlap:2006kj}
+ where $d$ is the `data', here the fiducial $\Lambda$ CDM correlation function in redshift and angular separation bins; $t$ is the alternative gravity model correlation function in those bins, $n_{\rm o}=75$ is the number of realisations of correlation functions used in the calculation of $C_{\rm cos}$ and $n_{\rm b}=48$ is the total number of bins in angular separation and redshift."
+ Note that we use the sample covariance estimate from Ilorizon (which follows CDM) for both ACDAL anc QCDAL cases: the QCDAL error bars should therefore only be considered as the correct order of magnitude., Note that we use the sample covariance estimate from Horizon (which follows $\Lambda$ CDM) for both $\Lambda$ CDM and QCDM cases; the QCDM error bars should therefore only be considered as the correct order of magnitude.
+ We calculate. for cach of our models the dillerence in 47 between the modified gravity mocel and a dark cnerev model. (either ACDAL or QCDAL). applying WALAP|SNe]BAO priors.," We calculate for each of our models the difference in $\chi^2$ between the modified gravity model and a dark energy model (either $\Lambda$ CDM or QCDM), applying WMAP+SNe+BAO priors."
+ Note that for ACDAL ancl f(R). we used the ACDAL backeround (Ixomatsuetal.2009). and for DGP and QCDAL the DOP background. 2008) was used (see 811).," Note that for $\Lambda$ CDM and $f(R)$, we used the $\Lambda$ CDM background \citep{Komatsu:2008hk} and for DGP and QCDM, the DGP background \citep{Fang:2008kc} was used (see 1)."
+ Figure 5. shows example results for our erouncd-based survey and Iuclid using the central cosmological parameter values for WALAP|SNe]BAO described in Sli: this is for the 2-1) projection case where we have not divided. the catalogue tomographicallv., Figure \ref{fig:correlation lcdm} shows example results for our ground-based survey and Euclid using the central cosmological parameter values for WMAP+SNe+BAO described in 1; this is for the 2-D projection case where we have not divided the catalogue tomographically.
+ We see from figures (a) and (b) that the cillerence between models is substantially greater in the nonlinear regime (0 309 than in the linear regime (0 309. as is the amplitude of the signal.," We see from figures (a) and (b) that the difference between models is substantially greater in the nonlinear regime $\theta \la 30'$ ) than in the linear regime $\theta \ga 30'$ ), as is the amplitude of the signal."
+ As (ο) and (d) show. it is also the case that the linear correlation function is small in the low-@ regime. if nonlinear corrections are not included.," As (c) and (d) show, it is also the case that the linear correlation function is small in the $\theta$ regime, if nonlinear corrections are not included."
+ We present the 7. dillorences between the modified eravities and fiducial dark energy models in Table 1.. for the 2-D (non-tomographic) cases including non-linear power.," We present the $\chi^2$ differences between the modified gravities and fiducial dark energy models in Table \ref{chi sq table}, for the 2-D (non-tomographic) cases including non-linear power."
+ We see that there is indeed strong discriminatory power between mocified gravity models. and ACDAL with the notional eround-based survey: the precision of Luclic is even more impressive., We see that there is indeed strong discriminatory power between modified gravity models and $\Lambda$ CDM with the notional ground-based survey; the precision of Euclid is even more impressive.
+ We also compare the constraints on DGaP ancl a QCDAL model of the same expansion history (ic. a DCP background)., We also compare the constraints on DGP and a QCDM model of the same expansion history (i.e. a DGP background).
+ Phe correlation functions for these models are shown in Figure 6.., The correlation functions for these models are shown in Figure \ref{fig:correlation qcdm}. .
+" One can either consider a QCDAL mocel with cosmological parameters equal to their central values inafit to WAIAP|BAOSNe. or more realistically the best fit QCDM model to the DCP mocel obtained by varying Qu, and ox."," One can either consider a QCDM model with cosmological parameters equal to their central values in a fit to WMAP+BAO+SNe, or more realistically the best fit QCDM model to the DGP model obtained by varying $\Omega_{\rm m}$ and $\sigma_8$."
+" We see that there is a choice of ο, and ox that make the QCDAL and DOP models. virtually indistinguishable.", We see that there is a choice of $\Omega_m$ and $\sigma_8$ that make the QCDM and DGP models virtually indistinguishable.
+ This is confirmed. by the bottom: row of ‘Table 1... which shows that the dillerenee in X7 for DGP and this QCDAL is insignificant.," This is confirmed by the bottom row of Table \ref{chi sq table}, which shows that the difference in $\chi^2$ for DGP and this QCDM is insignificant."
+ This is clearly partly due to the existence of a QCDM model with rather similar growth to the DOP. but also because of the low amplitude of the DCP correlation function. with the result that the error bars are larger in proportion to the signal than for other models.," This is clearly partly due to the existence of a QCDM model with rather similar growth to the DGP, but also because of the low amplitude of the DGP correlation function, with the result that the error bars are larger in proportion to the signal than for other models."
+ The power of future surveys to discriminate between eravitv models is borne out by the tomographic results., The power of future surveys to discriminate between gravity models is borne out by the tomographic results.
+ Examples of these are shown in Figure 7.. where we see the different redshift evolutions and amplitudes of the signal in the dillerent gravities.," Examples of these are shown in Figure \ref{fig:correlation different z}, where we see the different redshift evolutions and amplitudes of the signal in the different gravities."
+ Table 2/— confirms that. using the redshift) information allords us better discrimination between dark energy and mocified gravity models in every case. bv a factor of 50 to10054.," Table \ref{chi sq table 2} confirms that using the redshift information affords us better discrimination between dark energy and modified gravity models in every case, by a factor of 50 to."
+.. Because of this. we will only consider tomographic results from now on in the paper.," Because of this, we will only consider tomographic results from now on in the paper."
+ ‘Table 3 shows the impact of including non-linear power on our ability to discriminate between modified: &ravities., Table \ref{linear non-linear table} shows the impact of including non-linear power on our ability to discriminate between modified gravities.
+ Comparing these results with Table 2. we can see the improvement that measurements from the non-linear regime of the correlation functions. provide., Comparing these results with Table \ref{chi sq table 2} we can see the improvement that measurements from the non-linear regime of the correlation functions provide.
+ The improvement. is very substantial. amounting to an orderof magnitude in V dilference.," The improvement is very substantial, amounting to an orderof magnitude in $\chi^2$ difference."
+ ]t is important to note that. using only the Smith formula. without the CR asvmptote. causes," It is important to note that using only the \cite{Smith:2002dz} formula, without the GR asymptote, causes"
+"Because 5,<<39. the reeion to the left of the solid lines is excluded.","Because $\gamma_r\leq\gamma_0$, the region to the left of the solid lines is excluded."
+ In most cases. the dependence on the viewing auele à in the denominator of eqs. (23))(21))," In most cases, the dependence on the viewing angle $\alpha$ in the denominator of eqs. \ref{eq:star1}) \ref{eq:star2}) )"
+" cau be ueelected. lence pla.Nyy)7plix2.24.5,)sina0n. and similarly for pl (compare solid aud dotted lines in 39)."," can be neglected, hence $p^{\parallel}_\star(\alpha,\gamma_0,\gamma_r)\approx
+p^{\parallel}_\star(\pi/2,\gamma_0,\gamma_r)\,\sin^2\alpha$, and similarly for $p^\perp_\star$ (compare solid and dotted lines in )."
+" For the extreme case of perfect aliguiment between ay and J during both UW absorption (sy=x) aud IR Cluission (5,= o). corresponding to case (d) in retsec:estimate.. eqs. (23))"," For the extreme case of perfect alignment between $\ba$ and $\bJ$ during both UV absorption $\gamma_0=\infty$ ) and IR emission $\gamma_r=\infty$ ), corresponding to case (d) in \\ref{sec:estimate}, eqs. \ref{eq:star1}) )"
+ and (21)) may be simplified: The degree of polarization for in-plane aud out-of-plane modes from a population of PAID molecules above the center of a unifornrbriehtuess disk galaxw is obtained from(21): where Q=22(1coxw) is the angle subtended by the ealaxy as seen from. the cutting eraius., and \ref{eq:star2}) ) may be simplified: The degree of polarization for in-plane and out-of-plane modes from a population of PAH molecules above the center of a uniform-brightness disk galaxy is obtained from: where $\Omega=2\pi(1-\cos\omega)$ is the angle subtended by the galaxy as seen from the emitting grains.
+ The limit ο>0 recovers the case of a poiut-like ilhuuinatius source in eqs. (23)), The limit $\Omega\rightarrow0$ recovers the case of a point-like illuminating source in eqs. \ref{eq:star1}) )
+ aud (21)). whereas for au infinite disk (ο. 2x) the deeree of polarization tends to zero.," and \ref{eq:star2}) ), whereas for an infinite disk $\Omega\rightarrow2\pi$ ) the degree of polarization tends to zero."
+ Figure | shows. for the most favorable viewing geometry (0= w/2). the dependence of the degree of polarization on 50 for PATIs illuminated by a disk galaxy. with O= πι the different curves correspond to different values of 29/5. as explained in refsecistarpol. ," Figure \ref{fig:gal} shows, for the most favorable viewing geometry $\alpha=\pi/2$ ), the dependence of the degree of polarization on $\gamma_0$ for PAHs illuminated by a disk galaxy with $\Omega=\pi$ ; the different curves correspond to different values of $\gamma_0/\gamma_r$, as explained in \\ref{sec:starpol}. ."
+For diffuse ilbluniuation with €=xz. we find levels of polarization that are ~10€ of the values found for wni-cdirectional illuuination.," For diffuse illumination with $\Omega=\pi$, we find levels of polarization that are $\sim40\%$ of the values found for uni-directional illumination."
+" The extreme case of erains with perfect internal aliguinent (a,|J during both UV absorption and IR enussjon. case (d) iu refsec:estimate)) vieldsWehave obtained analytic foxiuulae for the degree of polarization expected for the 3.3. 6.2. 7.7. 8.6. 11.3. 12.7."," The extreme case of grains with perfect internal alignment $\ba\parallel\bJ$ during both UV absorption and IR emission, case (d) in \\ref{sec:estimate}) ) yieldsWehave obtained analytic formulae for the degree of polarization expected for the 3.3, 6.2, 7.7, 8.6, 11.3, 12.7,"
+"letter. the following cosmological parameters are assumed: Hl,—""kmsMpe *.Q,, 207. Q4—0.73.","letter, the following cosmological parameters are assumed; ${\rm H}_0=71\,{\rm km\,s}^{-1}\,{\rm
+Mpc}^{-1}$ , $\Omega_m=0.27$ , $\Omega_{\Lambda}=0.73$."
+ Dye(2008) provides a detailed: description of the SELL reconstruction method., \citet{dye08} provides a detailed description of the SFH reconstruction method.
+ The purpose of the brief. outline presented. here is both for completeness and to give specific details of the procedure we have used in this implementation., The purpose of the brief outline presented here is both for completeness and to give specific details of the procedure we have used in this implementation.
+ The method divides a galaxys history into discrete blocks of time., The method divides a galaxy's history into discrete blocks of time.
+ This results in a relatively low resolution SELL but one that does not adhere to a prescribed. (ie... potentially biased) parametric form.," This results in a relatively low resolution SFH, but one that does not adhere to a prescribed (i.e., potentially biased) parametric form."
+ Using a svnthetic library of simple stellar population (SSP) SIZDs. the fluxes resulting from a constant SER. normalised to one solar mass in each block. as measured in the observer frame across à range of filters. are calculated.," Using a synthetic library of simple stellar population (SSP) SEDs, the fluxes resulting from a constant SFR normalised to one solar mass in each block, as measured in the observer frame across a range of filters, are calculated."
+ Finding the contribution of Dux from each block in cach filter that best fits a set of observed Lluxes is a linear problem that can be solved. exactly., Finding the contribution of flux from each block in each filter that best fits a set of observed fluxes is a linear problem that can be solved exactly.
+ Ehe solution directly vields the galaxys SELL and stellar mass., The solution directly yields the galaxy's SFH and stellar mass.
+ Hn this letter. we have used the stellar SED libraries of Alaraston(2005) ancl. for comparison. Bruzual&Charlot (2003)...," In this letter, we have used the stellar SED libraries of \citet{maraston05} and, for comparison, \citet{bruzual03}. ."
+ Starting with a SSP SED. LV. of metallicity Z. a composite stellar population (CSP) SED. Li. is generated for the /th block of constant star formation in a given galaxy using dU Lyc)0) where the block spans the period. /;4 to ἐν in the ealaxv history and 7 is the age of the galaxy. (ie... the age of thes Universe today. minus the look-back time to the galaxy).," Starting with a SSP SED, $L_{\lambda}^{\rm SSP}$, of metallicity $Z$, a composite stellar population (CSP) SED, $L_{\lambda}^{\, i}$, is generated for the $i$ th block of constant star formation in a given galaxy using = t' (z)-t') where the block spans the period $t_{i-1}$ to $t_i$ in the galaxy's history and $\tau$ is the age of the galaxy (i.e., the age of the Universe today minus the look-back time to the galaxy)."
+" Phe quantity 2M;=f;¢;4, ensures that the CSP is normalised to one solar mass.", The quantity $\Delta t_i=t_i-t_{i-1}$ ensures that the CSP is normalised to one solar mass.
+ Ehe above integral is evaluated ον interpolating linearly in log(/) between the discrete time intervals at which the SSP SEDs are given in the libraries., The above integral is evaluated by interpolating linearly in $t$ ) between the discrete time intervals at which the SSP SEDs are given in the libraries.
+ ote that the method assumes that Z does not vary with age., Note that the method assumes that $Z$ does not vary with age.
+ Vo model the effects of extinction on the final SED ανα. he SED from all blocks in the SELL). reddening is applied.," To model the effects of extinction on the final SED (i.e., the SED from all blocks in the SFH), reddening is applied."
+" ‘This is achieved by individually reddening the CSP of each Mock using Lig=Lilo""2v/v whore dy is the extinction. Ay=4.05 and the Calzetti Law (Calzetti2000) is used for &(A)."," This is achieved by individually reddening the CSP of each block using $L_{\lambda,R}^{\, i} = L_{\lambda}^{\, i} \, 10^{-0.4
+k(\lambda)A_V/R_V}$ where, $A_V$ is the extinction, $R_V=4.05$ and the Calzetti Law \citep{calzetti00} is used for $k(\lambda)$ ."
+ Phe model Dux (ie. photon count) observed in filter j from a given. block 7 in the SELL when the ealaxy lies at à redshift z is then where d; is the luminosity distance and 7; is the transmission curve of filter j.," The model flux (i.e., photon count) observed in filter $j$ from a given block $i$ in the SFH when the galaxy lies at a redshift $z$ is then where $d_L$ is the luminosity distance and $T_j$ is the transmission curve of filter $j$."
+ Yo find the normalisations e; which result in a set of model fluxes that best fits the observed fluxes. the following V function is minimised where p is the galaxy [ux observed in filter j and σι is its error.," To find the normalisations $a_i$ which result in a set of model fluxes that best fits the observed fluxes, the following $\chi^2$ function is minimised where $F^{\rm obs}_j$ is the galaxy flux observed in filter $j$ and $\sigma_j$ is its error."
+ The sum in / acts over all Nia. SEL blocks., The sum in $i$ acts over all $N_{\rm block}$ SFH blocks.
+ The minimum. X47 occurs when the condition ὃν0 is simultaneously. satisfied. for alle;.," The minimum $\chi^2$ occurs when the condition $\partial\chi^2/\partial\,a_i=0$ is simultaneously satisfied for all$a_i$."
+ Phis gives a set of equations linear in the e; which are solvable using a standard matrix inversion (secDye2008.formorecletails)..," This gives a set of equations linear in the $a_i$ which are solvable using a standard matrix inversion \citep[see][for more
+details]{dye08}."
+" The a; are the stellar masses formed in cach block so that the total stellar mass of the galaxy is the sum AL,=Σολ5a;"," The $a_i$ are the stellar masses formed in each block so that the total stellar mass of the galaxy is the sum ${\rm M}_* =
+\sum_i^{N_{\rm block}} a_i$."
+ The SERs and hence the SELL is then given directly by dividing the e; by the time spannecl by cach corresponding block., The SFRs and hence the SFH is then given directly by dividing the $a_i$ by the time spanned by each corresponding block.
+ Formal errors on the e; are obtained. [rom the covariance matrix. computed in a simple additional step.," Formal errors on the $a_i$ are obtained from the covariance matrix, computed in a simple additional step."
+ To allow for uncertainty in source redshift. we performed a Monte Carlo simulation. randomising the redshift according to its error. and combined the resulting scatter in the e; in quadrature with the formal errors.," To allow for uncertainty in source redshift, we performed a Monte Carlo simulation, randomising the redshift according to its error, and combined the resulting scatter in the $a_i$ in quadrature with the formal errors."
+ The total error on the stellar mass was obtained in the same manner., The total error on the stellar mass was obtained in the same manner.
+ As discussed in Dye(2008)... regularisation must. be applied to ensure that the linear solution is well defined.," As discussed in \citet{dye08}, regularisation must be applied to ensure that the linear solution is well defined."
+ The strength of regularisation is controlled by a parameter referred to as the regularisation weight. denoted i hereafter.," The strength of regularisation is controlled by a parameter referred to as the regularisation weight, denoted $w$ hereafter."
+ The procedure outlined in the previous section is à single linear step which computes the SELL that gives the best fit (i.c.. minimuni v) to an observed set. of Dluxes for a given set of parameters 2. Z. y. Nia and ws," The procedure outlined in the previous section is a single linear step which computes the SFH that gives the best fit (i.e., minimum $\chi^2$ ) to an observed set of fluxes for a given set of parameters $z$, $Z$ , $A_V$, $N_{\rm block}$ and $w$."
+ This step is nested inside a non-linear search for the set of parameter values that gives the best overall fit to the observed [üxes., This step is nested inside a non-linear search for the set of parameter values that gives the best overall fit to the observed fluxes.
+ As discussed in Dye.(2008).. finding the best. &lobal solution can not be achieved by minimising X7 because the ellective number of degrees of freedom. depends on w in an unquantifiable manner.," As discussed in \citet{dye08}, finding the best global solution can not be achieved by minimising $\chi^2$ because the effective number of degrees of freedom depends on $w$ in an unquantifiable manner."
+ £t is therefore not. possible to use 47 to make a [air comparison of the goodness-of-fit between two parameter sets with cilfering values of «e., It is therefore not possible to use $\chi^2$ to make a fair comparison of the goodness-of-fit between two parameter sets with differing values of $w$.
+ To make a fair comparison. one must turn to Bavesian statistics and treat regularisation as a prior.," To make a fair comparison, one must turn to Bayesian statistics and treat regularisation as a prior."
+ In this wav. the Bayesian. evidence. denoted. € hereafter. allows dillerent. sets of parameters to be ranked fairly.," In this way, the Bayesian evidence, denoted $\epsilon$ hereafter, allows different sets of parameters to be ranked fairly."
+ The best global solution is that which maximises c., The best global solution is that which maximises $\epsilon$.
+ We investigated a range of cdillerent schemes to maximise c., We investigated a range of different schemes to maximise $\epsilon$.
+ The most cllicient anc reliable scheme that we found combines a standard &rid search with a downhill simplex minimisation to find the minimum of Inc (hence he maximum. of e)., The most efficient and reliable scheme that we found combines a standard grid search with a downhill simplex minimisation to find the minimum of $-\ln \epsilon$ (hence the maximum of $\epsilon$ ).
+ We use the simplex routine. linearly computing the SELL and evaluating € each time. to find the yest pair of values of zl andloga whilstkeeping Z and Nitecfixed.," We use the simplex routine, linearly computing the SFH and evaluating $\epsilon$ each time, to find the best pair of values of $A_V$ and$\log w$ whilstkeeping $Z$ and $N_{\rm block}$fixed."
+" Then. at the outer-most level. we step through a grid ofregularly spaced trial values of losZ and Ni; over heranges 2<log,,(Z/Z.)0.9 and2«Nia< 5."," Then, at the outer-most level, we step through a grid ofregularly spaced trial values of $\log Z$ and $N_{\rm block}$ over the ranges $ -2 <
+\log_{10} (Z/Z_\odot) < 0.3 $ and $2 < N_{\rm block} < 5$ ."
+ The value of z was fixed at the measured redshift of thegalaxy (sce Section 3)) at all times., The value of $z$ was fixed at the measured redshift of thegalaxy (see Section \ref{sec_data}) ) at all times.
+ Dyectal...(2000). found optical counterparts to 114 of the ~130 radio- and/or pam-identifiecBLAST sources located. in the region covered. by their optical data.," \citet{dye09} found optical counterparts to 114 of the $\sim 130$ radio- and/or $\,\mu$ m-identifiedBLAST sources located in the region covered by their optical data."
+ All, All
+The surface photometry for each galaxy. was derived (roi ELLIPSE fits to the final combiued üunages.,The surface photometry for each galaxy was derived from ELLIPSE fits to the final combined images.
+ Prior to ELLIPSE fitting. the foreground stars were masked aud repaced by the interpolated backgrouxd of neighboring pixels.," Prior to ELLIPSE fitting, the foreground stars were masked and replaced by the interpolated background of neighboring pixels."
+ The surface |lehtuess profiles. position angle. aud elipicity are show Las a function of semi-major axis in Fy[n]oOure L..," The surface brightness profiles, position angle, and ellipticity are shown as a function of semi-major axis in Figure \ref{fig:surf}."
+ Since the morphoogical paranetels were simil:w for all fiters. the apertures clefinec by 1re R-baud fit were used for subsequeu alalysis of the 5urlace briguuess profiles auc surface colors.," Since the morphological parameters were similar for all filters, the apertures defined by the R-band fit were used for subsequent analysis of the surface brightness profiles and surface colors."
+ The morphological parameters are sumiuaized in Table 2.., The morphological parameters are summarized in Table \ref{tab:results}.
+ Listed i1 Table 2 is 1je axis at wüch the B-aud surfacebrightsess drops to yn25 mag 37. Ros. aud the average ellipticity aud position angle (rom tle isoj»hota fits of the R-bancl images for ellipses wihi axes between anc Ro.," Listed in Table \ref{tab:results} is the semi-major axis at which the B-band surfacebrightness drops to 25 mag $^{-2}$, $_{25}$, and the average ellipticity and position angle from the isophotal fits of the R-band images for ellipses with semi-major axes between and $_{25}$."
+ Apertu‘es larger than were used for these average valtes slice several of the galaxies in tus sample have brieαἱ nuclear regions., Apertures larger than were used for these average values since several of the galaxies in this sample have bright nuclear regions.
+ Also listed in Table 2. are B-baud apparent magnitudes at which tle curves of erowtli level off aud the observe optical colors as measured within the Res; aperture., Also listed in Table \ref{tab:results} are the B-band apparent magnitudes at which the curves of growth level off and the observed optical colors as measured within the $_{25}$ aperture.
+ As illustrated in Figure .. the surface colors auc apert colors reaci ΜΜ values near Ros so tjese values were chosen as the most represeative gal:Wy color.," As illustrated in Figure \ref{fig:surf}, the surface colors and aperture colors reach similar values near $_{25}$ so these values were chosen as the most representative galaxy color."
+ Also tabulated in Table 2 are est]mates of the foreground Galactic exjuctioi as calculaed from the dist maps of Schlegeletal.(1998).," Also tabulated in Table \ref{tab:results}
+ are estimates of the foreground Galactic extinction as calculated from the dust maps of \citet{SFD98}."
+ The iilages were examiued for residual disky or boxy isophotes by stXracting the model surface photometry from the image., The images were examined for residual disky or boxy isophotes by subtracting the model surface photometry from the image.
+ Eleven of the 16 galaxies had residua] structure hat was readily apparent in the residual images: the presence of cdisky or boxy isopoles is tabtlated in Table 2.., Eleven of the 16 galaxies had residual structure that was readily apparent in the residual images; the presence of disky or boxy isophotes is tabulated in Table \ref{tab:results}.
+" Deviatious from the best-fittiug ellipse as measures by B, are show in Figure 1..", Deviations from the best-fitting ellipse as measured by $_4$ are shown in Figure \ref{fig:surf}.
+ As p'eviously noted by Celiaetal.(2003).. there appeas to be little cor‘elation between the sresence of clisky or boxy residuals and the kinematic prope‘ties of the £alaxies.," As previously noted by \citet{GGvM03}, there appears to be little correlation between the presence of disky or boxy residuals and the kinematic properties of the galaxies."
+ As ab ildication ol the kinematic properties of these galaxies. the rotation cirve slope f.'on vanZeeetal.(2001) is listed i1 Table 2..," As an indication of the kinematic properties of these galaxies, the rotation curve slope from \citet{vSH04} is listed in Table \ref{tab:results}."
+ Due to the limited spatial extent of the stellar kinenatic Observations. we list he rotation curve slope rather than a iaxiumunm velocity widhi: galaxies wit1 rotation curve slopes ereater tlau 20 kus ! | have significant rotation.," Due to the limited spatial extent of the stellar kinematic observations, we list the rotation curve slope rather than a maximum velocity width; galaxies with rotation curve slopes greater than 20 km $^{-1}$ $^{-1}$ have significant rotation."
+" Based on the kieludic observations. VCC Ίτο, 137. 513. 965. 990. 1036. 1713. 18257. aud 2019 form the subset of ‘olaion dominated cEs in this sample."," Based on the kinematic observations, VCC 178, 437, 543, 965, 990, 1036, 1743, 1857, and 2019 form the subset of rotation dominated dEs in this sample."
+ Table 3 summarizes the derived photometric properties of the dwa ‘Celiptical galaxies 1 this sample., Table \ref{tab:results2} summarizes the derived photometric properties of the dwarf elliptical galaxies in this sample.
+ The B-band absolute magnitude was calculated from the appareu magnitude assining all galaxies are at the Virgo Cluster distance of 16.1 Mpe (Ixelsonetal.2000) and. correcting Lor Galactic extinction., The B-band absolute magnitude was calculated from the apparent magnitude assuming all galaxies are at the Virgo Cluster distance of 16.1 Mpc \citep{Ke00} and correcting for Galactic extinction.
+ Also listed are the effective radius ancl ellective surface brightness in B-band., Also listed are the effective radius and effective surface brightness in B-band.
+ For all of the following except the exponential parameters. size parameters are quoted in geometric mean radii where r = Yab: for the exponential parameters. the semi-major axis is adopted as the radius.," For all of the following except the exponential parameters, size parameters are quoted in geometric mean radii where r = $\sqrt{ab}$ ; for the exponential parameters, the semi-major axis is adopted as the radius."
+ The effective radius tabulated in Table 3) is the halClight radius based ou the apparent, The effective radius tabulated in Table \ref{tab:results2} is the half-light radius based on the apparent
+sparse representation as described in ?..,sparse representation as described in \citet{Candes:2006p23}.
+ In this paper. we adopt the Daubechies D8 wavelet transforms (?) as the dictionary.," In this paper, we adopt the Daubechies D8 wavelet transforms \citep{daubechies1992tlw} as the dictionary."
+ Other wavelet transforms can also be adopted. the selection depends on the property of the Faraday dispersion function.," Other wavelet transforms can also be adopted, the selection depends on the property of the Faraday dispersion function."
+ We choose the D8 wavelet transform. because we assume that the Faraday dispersion function with thick sources is a sine-like signal.," We choose the D8 wavelet transform, because we assume that the Faraday dispersion function with thick sources is a sinc-like signal."
+ We can rewrite Eq. (8)), We can rewrite Eq. \ref{explainRM}) )
+ as wwhere W! is the inverse wavelet transform matrix of size Nx Nia as the wavelet coefficient of the Faraday dispersio function f., as where $\mathrm{W}^{-1}$ is the inverse wavelet transform matrix of size $N \times N$; $\alpha$ is the wavelet coefficient of the Faraday dispersion function $\mathbf{f}$.
+ The wavelet transform matrix is denoted as W. therefore. Y=Wf.," The wavelet transform matrix is denoted as $\mathrm{W}$, therefore, $\alpha=\mathrm{W}\mathbf{f}$."
+ Other symbols follow the definitions 1 Eq. (8)., Other symbols follow the definitions in Eq. \ref{explainRM}) ).
+ Under the condition that Yf=p. we adopt th0 following assumption: both the real part and the imaginary para of the Faraday thick sources will have a sparse representation 1 the wavelet domain independently.," Under the condition that $\mathrm{Y}\mathbf{f}=\widetilde{\mathbf{p}} $, we adopt the following assumption: both the real part and the imaginary part of the Faraday thick sources will have a sparse representation in the wavelet domain independently."
+ Then the wavelet based CS RM synthesis nethod for Faraday thick sources can be writte as This CS-based rotation measure synthesis for Faraday thick sources is abbreviated as CS-RM-Thick., Then the wavelet based CS RM synthesis method for Faraday thick sources can be written as This CS-based rotation measure synthesis for Faraday thick sources is abbreviated as CS-RM-Thick.
+ So far. we have proposed two RM synthesis methods: CS-RM-Thin and CS-RM-Thick for solving Faraday thin sources and thick sources. respectively.," So far, we have proposed two RM synthesis methods: CS-RM-Thin and CS-RM-Thick for solving Faraday thin sources and thick sources, respectively."
+ However. this begs the question: which method should be selected if there are both Faraday thin sources and thick sources along the line of sight?," However, this begs the question: which method should be selected if there are both Faraday thin sources and thick sources along the line of sight?"
+ Moreover. how can we make a selection if we have no prior information about the Faraday dispersion function. 1.8. we are not sure what it looks like?," Moreover, how can we make a selection if we have no prior information about the Faraday dispersion function, i.e. we are not sure what it looks like?"
+ Clearly. neither of them is suitable. we therefore need another solution for solving the above problems.," Clearly, neither of them is suitable, we therefore need another solution for solving the above problems."
+ Let us assume that there are both Faraday thin sources and Faraday thick sources in F(6)., Let us assume that there are both Faraday thin sources and Faraday thick sources in $F(\phi)$ .
+" Suppose fi, denotes the Faraday thin sources in. F(@) in a vector format of length N: fiji, denotes the Faraday thick sources in F(@) in a vector format of length N. then fii;+fini,=£f. Eq.(8))"," Suppose $\mathbf{f}_{\rm{thin}}$ denotes the Faraday thin sources in $F(\phi)$ in a vector format of length $N$; $\mathbf{f}_{\rm{thick}}$ denotes the Faraday thick sources in $F(\phi)$ in a vector format of length $N$, then $\mathbf{f}_{\rm{thin}}+\mathbf{f}_{\rm{thick}}=\mathbf{f}$. \ref{explainRM}) )"
+" can be rewritten as Since we know that the LI norm can preserve sparsity; Faraday thin sources show sparsity in the Faraday depth domain: Faraday thick sources show sparsity in the wavelet domain. we propose the following solution for the mixed circumstance where the definition of W is the same as the above subsection,"," can be rewritten as Since we know that the L1 norm can preserve sparsity; Faraday thin sources show sparsity in the Faraday depth domain; Faraday thick sources show sparsity in the wavelet domain, we propose the following solution for the mixed circumstance where the definition of $\mathrm{W}$ is the same as the above subsection."
+ The above solution for Faraday mixed sources ts still based on the spirit of CS by preserving the sparsity in the Faraday depth domain for Faraday thin sources and in the wavelet domain for Faraday thick sources. simultaneously.," The above solution for Faraday mixed sources is still based on the spirit of CS by preserving the sparsity in the Faraday depth domain for Faraday thin sources and in the wavelet domain for Faraday thick sources, simultaneously."
+ This CS-based rotation measure synthesis for Faraday mixed sources is abbreviated as CS-RM-Mix., This CS-based rotation measure synthesis for Faraday mixed sources is abbreviated as CS-RM-Mix.
+ In this section. the implementation details of the above three proposed CS-based RM synthesis methods will be given.," In this section, the implementation details of the above three proposed CS-based RM synthesis methods will be given."
+ To create à general experiment layout. we borrow some definitions and conclusions from ?..," To create a general experiment layout, we borrow some definitions and conclusions from \cite{Brentjens:2005p3385}."
+ Both the diagram of the wavelength square οἱ and # are displayed in figure 1.., Both the diagram of the wavelength square $\lambda^2$ and $\phi$ are displayed in figure \ref{fig_wavelength}.
+ The maximum observable Faraday depth [όν] is given in ? wwhere Ot Dis the width of an observing channel.," The maximum observable Faraday depth $\| \phi_{\mathrm{max}}\|$ is given in \cite{Brentjens:2005p3385}
+ where $\delta \lambda^2$ is the width of an observing channel."
+ The full width at half maximum (FWHM) of the main peak of the RMSF can be estimated by where A;U is the width of the total P distribution., The full width at half maximum (FWHM) of the main peak of the RMSF can be estimated by where $\bigtriangleup\lambda^2$ is the width of the total $\lambda^2$ distribution.
+ Before using CS-based RM synthesis. the following two steps are needed: The selection of these CS-based RM synthesis methods depends on the prior knowledge about the Faraday dispersion function.," Before using CS-based RM synthesis, the following two steps are needed: The selection of these CS-based RM synthesis methods depends on the prior knowledge about the Faraday dispersion function."
+ If we assume 1t includes Faraday thin sources onlyalong the line of sight. we should select CS-RM-Thin.," If we assume it includes Faraday thin sources onlyalong the line of sight, we should select CS-RM-Thin."
+ On the other hand. CS-RM-Thick should be used if we know that," On the other hand, CS-RM-Thick should be used if we know that"
+ (fae).,$f_{\rm det}$ ).
+ Por the idealised photon-noise-limitecd case objects as small as Mercury. could be detected to periods of around dd. and the Moon for periods less than Edd. Once the impact correlated instrumental noise (red noise) is added. Aloon-sized companions become almost undetectable. though the recovery| rates for larger bodies. particularly in short-period orbits. remains encouraging.," For the idealised photon-noise-limited case objects as small as Mercury could be detected to periods of around d, and the Moon for periods less than d. Once the impact correlated instrumental noise (red noise) is added, Moon-sized companions become almost undetectable, though the recovery rates for larger bodies, particularly in short-period orbits, remains encouraging."
+ Our key conclusion from these simulations is that for the case of transits of WDs. the degree of photometric precision delivered by a survey is of somewhat secondary importance compared to a high cadence ancl continuous coverage.," Our key conclusion from these simulations is that for the case of transits of WDs, the degree of photometric precision delivered by a survey is of somewhat secondary importance compared to a high cadence and continuous coverage."
+ For planet-sized bodies individual transits will be quite deep aud readily detectable in data of moderate photometric quality. however it is the short duration of the transits that is the main factor limiting the transit detection rate in surveys optimised [or main sequence stars.," For planet-sized bodies individual transits will be quite deep and readily detectable in data of moderate photometric quality, however it is the short duration of the transits that is the main factor limiting the transit detection rate in surveys optimised for main sequence stars."
+ Encouraged by the results of our simulations we selected a sample of WDs. which have been routinely. monitored by WASP through the 20042008 observing seasons. ancl performed. a systematic search. for eclipsing and transiting substellar ancl planctary companions.," Encouraged by the results of our simulations we selected a sample of WDs, which have been routinely monitored by WASP through the 2004–2008 observing seasons, and performed a systematic search for eclipsing and transiting substellar and planetary companions."
+ We selected: the sample by cross-correlating the catalogue of WASP objects for which more than 600 data points are available with the AleCook & Sion catalogue (?).., We selected the sample by cross-correlating the catalogue of WASP objects for which more than 600 data points are available with the McCook $\&$ Sion catalogue \citep{WDcat}.
+ Phe resulting sample of 194 WDs with magnitude V«15 is presented in Table 4., The resulting sample of 194 WDs with magnitude $\textrm{V}<15$ is presented in Table 4.
+ We searched. the sample for transits and eclipses using our implementation of the BLS algorithim. searching periods ranging from 2 hours to 15 days.," We searched the sample for transits and eclipses using our implementation of the BLS algorithm, searching periods ranging from 2 hours to 15 days."
+ In addition we have also inspected. each of the individual light curves by eve., In addition we have also inspected each of the individual light curves by eye.
+ In both searches we found no evidence for any transiting ancl eclipsing companions within the period. range searched. in this study., In both searches we found no evidence for any transiting and eclipsing companions within the period range searched in this study.
+ We have used this null result together with the results of our simulations to estimate an upper-Iinit to the frequeney of such close companions for the sample of WDs considered in this study., We have used this null result together with the results of our simulations to estimate an upper-limit to the frequency of such close companions for the sample of WDs considered in this study.
+ In order to estimate an upper limit to the frequency of close substellar ancl planetary. companions to WDs. we used the detection limits derived (rom our simulations and the results obtained. from the analysis of the sample of 194 WDs.," In order to estimate an upper limit to the frequency of close substellar and planetary companions to WDs, we used the detection limits derived from our simulations and the results obtained from the analysis of the sample of 194 WDs."
+ We first used a binomial distribution to describe the probability, We first used a binomial distribution to describe the probability
+As cliscussecl in the Lut‘odduction. oue of the novelties of the present. work is the use. [or some oL the most imiportant. chetical processes. of rates obtained [rom accurate qlautun calculations aud not rou qualitative estünates.,"As discussed in the Introduction, one of the novelties of the present work is the use, for some of the most important chemical processes, of rates obtained from accurate quantum calculations and not from qualitative estimates."
+ It is therefore important for the reactions Ivoving o recall the beh:wior of the al-iuitio reactive rates (in the temperature raiee ol interest) tliat lave bee obtained via the quantum iiethods outlined in Section 2., It is therefore important for the reactions involving $^+$ to recall the behavior of the ab-initio reactive rates (in the temperature range of interest) that have been obtained via the quantum methods outlined in Section 2.
+ The data 'eported in Figure 1 iudicate he depletion rates or the cis:ippearance of LiH molecules by exothlierijcally reacting with he surrounding atomic liverogen (toj» panel)., The data reported in Figure 1 indicate the depletion rates for the disappearance of LiH molecules by exothermically reacting with the surrounding atomic hydrogen (top panel).
+ The temperature behavior of he ab-initio data is Clearly very differeut. [rom that assumed by SLD96., The temperature behavior of the ab-initio data is clearly very different from that assumed by SLD96.
+ The new calculations of Bovi10 et al. (, The new calculations of Bovino et al. (
+2009) inclicate a variatiou over the same raiee of nearly three orders of magnitude.,2009) indicate a variation over the same range of nearly three orders of magnitude.
+ T1ο same occurs for he destruction reaction of LiH by inte‘action with protons (mniddle panel): the ab-initio calculations (Bovino et al., The same occurs for the destruction reaction of LiH by interaction with protons (middle panel): the ab-initio calculations (Bovino et al.
+" 2010b) show that the ""ate varies by about one order of maguitide over the same ange of temperatures.", 2010b) show that the rate varies by about one order of magnitude over the same range of temperatures.
+ Finally. the €uantuni rates associated with the clestrtclion of by he surrouncliie hydrogeu gas (bottom panel) couliriu the weak temperature depeucence reported earlier (Lepp et al.," Finally, the quantum rates associated with the destruction of $^+$ by the surrounding hydrogen gas (bottom panel) confirm the weak temperature dependence reported earlier (Lepp et al."
+ 2002). although they turu out to be about a factor of E larger than tlie estimate of SLD96.," 2002), although they turn out to be about a factor of 4 larger than the estimate of SLD96."
+ Ouce he new quautities are eniployed within the evolutionary scheme ouliued in Section 3. they vield a new set of production/destruction rates of LiH and as a fueion of redshift. as shown in Fietre 2 ancl 3.," Once the new quantities are employed within the evolutionary scheme outlined in Section 3, they yield a new set of production/destruction rates of LiH and $^+$ as a function of redshift, as shown in Figure 2 and 3."
+") The various curves shown in Fieuree 2 follow the notation [weeiven in Table 1 to the most relevant [m'Ocesses: those labelled as py...-.py Golid lines) reer to the different reactious leadiug to LiH [m'oductiou. while those labelled dy...d, (clashed lines) refer to the dest""uction processes."," The various curves shown in Figure 2 follow the notation given in Table 1 to the most relevant processes: those labelled as $p_1, \ldots,
+p_4$ (solid lines) refer to the different reactions leading to LiH production, while those labelled $d_1, \ldots, d_4$ (dashed lines) refer to the destruction processes."
+ A look atte figure allows oue to make the folowing cotells: (7)) The production by radiative recombination Iurom the Li(?P) stae (pi) clearly dominates ie redshifts. while the procluction by recombin:ion from the Li(-SE ) state (po) becomes the dominant route [|‘on 2£SOO down ο 2~NO.," A look at the figure allows one to make the following comments: ) The production by radiative recombination from the $^2P$ ) state $p_1$ ) clearly dominates at high redshifts, while the production by recombination from the $^2S$ ) state $p_2$ ) becomes the dominant route from $z\simeq 800$ down to $z\simeq 80$."
+ Then. the channel of 'ecombination driven by elecron attachment (pj) takes over.," Then, the channel of recombination driven by electron attachment $p_4$ ) takes over."
+ However. we mus note that the latter rate is not known [rom quautum caleulatious. but only estimated by SLD96. together with the esimates of LiH formation from (SLD96). (H)," However, we must note that the latter rate is not known from quantum calculations, but only estimated by SLD96, together with the estimates of LiH formation from $^-$ (SLD96). )"
+) At the|ighest recshifs. the clestruction chaiuel is dominated by photodissociatiou (di). a rate estimated rom the detailed balance of process (pa). while at lower + values (25CEN 300) the collisional process of reactive destruction ciscussed earier (do) takes over aud dominates the whole destruction clianiel down to tle lowest redshifts.," At the highest redshifts, the destruction channel is dominated by photodissociation $d_1$ ), a rate estimated from the detailed balance of process $p_2$ ), while at lower $z$ values $z\lesssim 300$ ) the collisional process of reactive destruction discussed earlier $d_2$ ) takes over and dominates the whole destruction channel down to the lowest redshifts."
+ The same reactive process with protous (da) is seen to be less important. although larger than previois estimates (dj).," The same reactive process with protons $d_3$ ) is seen to be less important, although larger than previous estimates $d_4$ )."
+ Iun conclusion. the above daa indicate that the raciative photoassociation from the Ξ25 level," In conclusion, the above data indicate that the radiative photoassociation from the $n=25$ level"
+surprisingly. it has received considerable attention from theoreticians. and in many wavs has been a paradigm lor gravitational interactions in multiplanet extrasolar planetary. systenis.,"surprisingly, it has received considerable attention from theoreticians, and in many ways has been a paradigm for gravitational interactions in multiplanet extrasolar planetary systems."
+ Al first. research focused on its stability Laughlin Adams 1999; Rivera Lissaner 2000: Lissaner Rivera 2001: Laskar 2000: Barnes Quinn 2001: Gozddziewski 2001).," At first, research focused on its stability Laughlin Adams 1999; Rivera Lissauer 2000; Lissauer Rivera 2001; Laskar 2000; Barnes Quinn 2001; Goźddziewski 2001)."
+ These investigations showed (he svstem appeared io lie near the edge of instability. although an additional body could survive in between planets b and c (Rivera Lissauer 2001).," These investigations showed the system appeared to lie near the edge of instability, although an additional body could survive in between planets b and c (Rivera Lissauer 2001)."
+ Later. attention turned to the apsidal motion Slepinski 2000: Chiang Murray 2002: Malhotra 2002: Ford 2005: Barnes Greenberg 200G6a.¢. 20072). with most investigators considering how the apsidal behavior provides clues to formation mechanisms such as planet-planet scattering or migration via disk torques.," Later, attention turned to the apsidal motion Stepinski 2000; Chiang Murray 2002; Malhotra 2002; Ford 2005; Barnes Greenberg 2006a,c, 2007a), with most investigators considering how the apsidal behavior provides clues to formation mechanisms such as planet-planet scattering or migration via disk torques."
+ The recent direct measurement of (he actual masses of planets e and d. and especially their 30° mutual inclination. via astrometry (AleArthur 2010) maintains (his svstenirs prominence among known exoplanetary svstenis.," The recent direct measurement of the actual masses of planets c and d, and especially their $^\circ$ mutual inclination, via astrometry (McArthur 2010) maintains this system's prominence among known exoplanetary systems."
+ In (his investigation. we evaluate planets c and d's gravitational interactions (ignoring b as its orbit is still only constrained bv radial velocity (RV) observations). as presented in MeArthur (2010). which place strict constraints on the svstem's origin.," In this investigation, we evaluate planets c and d's gravitational interactions (ignoring b as its orbit is still only constrained by radial velocity (RV) observations), as presented in McArthur (2010), which place strict constraints on the system's origin."
+ such large mutual inclinations among planets are unknown in our Solar System. ancl hence indicate different processes occurred during or after the planet formation process., Such large mutual inclinations among planets are unknown in our Solar System and hence indicate different processes occurred during or after the planet formation process.
+ We assume e and d formed in coplanar. low eccentricity orbits in a standard planetary formation models (see IIubiekvj 2010: Maver 2010). but. then. additional phenomena. occurring late or after the lormation process. altered the svstem’s architecture.," We assume c and d formed in coplanar, low eccentricity orbits in a standard planetary formation models (see Hubickyj 2010; Mayer 2010), but then additional phenomena, occurring late or after the formation process, altered the system's architecture."
+" We consider {wo plausible mechanisms: a) If the distant stellar companion D (Lowrance 2002) is on a significantly inclined (relative to the initial orbital plane of the planets) aud/or eccentric orbit. i may pump up eccentricities and inclinations (through ""Ixozai interactions (Ixozai 1962: Takeda 2007): or b) If the planets Form close together they may interact and scatter into mutually inclined. orbits. as shown in previous studies (Weidenschilling Marzari 2002: Chatterjee 2008: Ravimoncl 2010)."," We consider two plausible mechanisms: a) If the distant stellar companion B (Lowrance 2002) is on a significantly inclined (relative to the initial orbital plane of the planets) and/or eccentric orbit, it may pump up eccentricities and inclinations through “Kozai” interactions (Kozai 1962; Takeda 2007); or b) If the planets form close together they may interact and scatter into mutually inclined orbits, as shown in previous studies (Weidenschilling Marzari 2002; Chatterjee 2008; Raymond 2010)."
+ The mutual inclinations are obviously of the most interest. but other features of the svstem are also important.," The mutual inclinations are obviously of the most interest, but other features of the system are also important."
+ As represents the first exoplanetary svstem with [ull (hree-climensional orbits ancl (rue masses «του measured (aside from the pulsar svsteni PSR 1257412 [Wolszezan 1994]). we may exploit this information when evaluating formation models.," As represents the first exoplanetary system with full three-dimensional orbits and true masses directly measured (aside from the pulsar system PSR 1257+12 [Wolszczan 1994]), we may exploit this information when evaluating formation models."
+ To that end. we develop a simple metric that quantifies the success of a mocel αἱ reproducing numerous observed aspects of (he ο And planetary svstem.," To that end, we develop a simple metric that quantifies the success of a model at reproducing numerous observed aspects of the $\upsilon$ And planetary system."
+ Although we apply this approach to vc And. it is generalizable to anv planetary svstem.," Although we apply this approach to $\upsilon$ And, it is generalizable to any planetary system."
+easily. stripped. than racial orbits: while racial orbits are more casily stripped. than retrograde orbits (Reacl 2006).,easily stripped than radial orbits; while radial orbits are more easily stripped than retrograde orbits \cite{read06}.
+. In order to understand the elfects of a tidal field on he particles (ic. particles of dark matter or stars?) within he satellite with mass AZ. we assume a satellite which experiences a constant tidal force (e.g. on a circular orbit within a spherical host potential with negligible dynamical riction).," In order to understand the effects of a tidal field on the particles (i.e. particles of dark matter or ) within the satellite with mass $M_{\rm sat}$, we assume a satellite which experiences a constant tidal force (e.g. on a circular orbit within a spherical host potential with negligible dynamical friction)."
+ uncealistAlthough the assumption of a constant tidal force is ic (Lew satellites will be on orbits that are close to circular and. for our default satellite. the dynamical friction imescale is of order I0Civr. resulting in significant change in orbital radius on a timescale comparable to the mass oss timescale from the satellite). we choose to explore this idealized case here ancl will explore more realistic examples in future work.," Although the assumption of a constant tidal force is unrealistic (few satellites will be on orbits that are close to circular and, for our default satellite, the dynamical friction timescale is of order 10Gyr, resulting in significant change in orbital radius on a timescale comparable to the mass loss timescale from the satellite), we choose to explore this idealized case here and will explore more realistic examples in future work."
+ Particle orbits within the satellite however. have arbitrary ellipticitv.," Particle orbits within the satellite however, have arbitrary ellipticity."
+ Figure 2. shows the geometry of the particle orbit considered. below., Figure \ref{fig:diag} shows the geometry of the particle orbit considered below.
+ The energy per unit mass of a particle is where e(r) is the gravitational potential of the satellite and e ds the velocity of the particle., The energy per unit mass of a particle is where ${\Phi}(x)$ is the gravitational potential of the satellite and $v$ is the velocity of the particle.
+ We assume an NEW potential for the satellite. resulting in a potential: where p is the mean density of the halo.," We assume an NFW potential for the satellite, resulting in a potential: where $\bar{\rho}$ is the mean density of the halo."
+ We consicer a rotating frame of reference which rotates ab the same rate as the orbitting satellite., We consider a rotating frame of reference which rotates at the same rate as the orbitting satellite.
+ As the particle moves around its orbit it experiences a torque (relative to the centre of mass of the satellite) due to tidal (from the host potential) and Coriolis forces., As the particle moves around its orbit it experiences a torque (relative to the centre of mass of the satellite) due to tidal (from the host potential) and Coriolis forces.
+ The angular momentum of the particle therefore varies around the orbit., The angular momentum of the particle therefore varies around the orbit.
+ The particle must always remain bound to the combined. satellite plus host potential (since those potentials are not time-varving the energy. of the particle is conserved. along its orbit)., The particle must always remain bound to the combined satellite plus host potential (since those potentials are not time-varying the energy of the particle is conserved along its orbit).
+ Llowever. we can ask the question: Lf the host potential were instantancously removed. would the particle be bound to the satellite or not?," However, we can ask the question: If the host potential were instantaneously removed, would the particle be bound to the satellite or not?"
+ Το answer this question we simply. need to know the angular momentum of the particle relative to the satellite centre. as this immediately gives the kinetic energy of the particle relative to the satellite centre.," To answer this question we simply need to know the angular momentum of the particle relative to the satellite centre, as this immediately gives the kinetic energy of the particle relative to the satellite centre."
+ We expect this criterion to be a reasonable discriminator between particles which are essentially. confined to the satellite and those which can move more freely throughout the host potential (i.e. those in tidal tails)., We expect this criterion to be a reasonable discriminator between particles which are essentially confined to the satellite and those which can move more freely throughout the host potential (i.e. those in tidal tails).
+ Experiments with N-bocky simulations (see refsec:nbods)) confirm this expectation., Experiments with N-body simulations (see \\ref{sec:nbody}) ) confirm this expectation.
+ The calculation. of the perturbation to the particles angular momentum is given in Appendix X... where we compute the maximuni value of the particles energy. (in a frame instantancously moving with the satellite). Duas.," The calculation of the perturbation to the particle's angular momentum is given in Appendix \ref{app:AM}, where we compute the maximum value of the particle's energy (in a frame instantaneously moving with the satellite), $E_{\rm max}$."
+ The total mass loss sullered [rom the satellite galaxy is found by integrating5 over phase space For this work we assume that the phase space density. F(x.v). is à separable function. so it can be written as: where p(.r) the profile of the satellite and ii(s)= lis) and. DNf(s)--=0 if s<0.," The total mass loss suffered from the satellite galaxy is found by integrating over phase space For this work we assume that the phase space density, $F(\bf{x,v})$ , is a separable function, so it can be written as: where $\rho(x)$is the density profile of the satellite and $H(s)=1$ if $s>0$ and $H(s)=0$ if $s<0$."
+ Hore. flr.ὁ) is the velocity PMion of the particles in the satellite for which we assume an (anisotropic in. general) Caussian and which is normalised. such that hodFledistribution»uxcsinodede=1.," Here, $f(v,\phi)$ is the velocity distribution function of the particles in the satellite for which we assume an (anisotropic in general) Gaussian distribution and which is normalised such that $\int_0^\infty
+\int_0^\pi f(v,\phi) 2 \pi v^2 \sin\phi \d\phi \d v=1$."
+ The radial density prolile can be foundin à similar way asshown in eqn. LO .," The radial density profile can be found in a similar way asshown in eqn. \ref{eq:mloss}, ,"
+ by simply not performing the integral over radius., by simply not performing the integral over radius.
+For the 21 jiu MIPS photometry we used the mosaics originally generated with the DAT (Cordonet and published bySicilia-Aguilaretal.(2006).. Teixciraetal.(2006).. aud Balogetal.(2007).,"For the 24 $\mu$ m MIPS photometry we used the mosaics originally generated with the DAT \citep{Gordon05} and published by\citet{Sici06}, \citet{Teix06}, and \citet{Balo07}."
+". First we measured the flux of the ""head via PSF fitting usine then we estimated the flux in the tail.", First we measured the flux of the “head” via PSF fitting using then we estimated the flux in the tail.
+ To do so. we selected a rectangular region along the tail iu the PSF subtracted muiage aud measured the signal im this region.," To do so, we selected a rectangular region along the tail in the PSF subtracted image and measured the signal in this region."
+ Next we moved the selected region to nearby backerounud positions on both sides of the tail and averaged the signal measured at these positions., Next we moved the selected region to nearby background positions on both sides of the tail and averaged the signal measured at these positions.
+ We then subtracted the averaged background from the sienal at the position of the tail to get the total fiux in the tail., We then subtracted the averaged background from the signal at the position of the tail to get the total flux in the tail.
+ The results of our photometry are plotted iu Figs. 2...3.. ," The results of our photometry are plotted in Figs. \ref{fig:1396IRSSL}, \ref{fig:2264IRSSL},"
+aud E. (black dots: head. orange dots: tail).," and \ref{fig:2244IRSSL} (black dots: head, orange dots: tail)."
+ We detected uo Pao. emission in uv of the objects: we will now derive upper lits to the amount of gas conrpatible with this result., We detected no $\alpha$ emission in any of the objects; we will now derive upper limits to the amount of gas compatible with this result.
+ Consider a spherical (or hemispherical) cloud of hydrogen with a radius r. located at a distance d from an Ο star that has a Lyinan coutimmun Duuinositv Q ün 5s ij," Consider a spherical (or hemispherical) cloud of hydrogen with a radius $r$, located at a distance $d$ from an O star that has a Lyman continuum luminosity $Q$ (in ${\rm \gamma s^{-1}}$ )."
+ Tf the cloud is sufficieutly dense. it will be optically thick to the Lxiuau continua radiation.," If the cloud is sufficiently dense, it will be optically thick to the Lyman continuum radiation."
+ Iu the model developed for the Orion. Nebula. proplyds (Johnstoneetal.1998).. soft. nondonizius UV heats the outer lavers ofthe cloud resulting in a quasi-steady-state. soft-UV. driven wind expaudiug with a velocity of a few kilometers per second.," In the model developed for the Orion Nebula proplyds \citep{John98}, soft, non-ionizing UV heats the outer layers of the cloud resulting in a quasi-steady-state, soft-UV driven wind expanding with a velocity of a few kilometers per second."
+ Ionizing Lyman continua radiation forms a quasi-stationary ionization frout (1-ront) iu this wind where the eas becomes ionized aud rveated to about 101 K. The resulting pressure jump diives a D-type shock iuto the expaucding low-velocity jeutral wind that compresses aud decelerates the flow fore if passes through the I-frout., Ionizing Lyman continuum radiation forms a quasi-stationary ionization front (I-front) in this wind where the gas becomes ionized and heated to about $10^4$ K. The resulting pressure jump drives a D-type shock into the expanding low-velocity neutral wind that compresses and decelerates the flow before it passes through the I-front.
+ At the Lfrout. the asma accelerates to about the sound speed iu the jonized eas. cry. and the pressure eradicut results in a spherically divergeut flow.," At the I-front, the plasma accelerates to about the sound speed in the ionized gas, $_{II}$, and the pressure gradient results in a spherically divergent flow."
+ Once ionization equilibria as been established in the iouized flow. aud the svstems Is in a steady state of photo-ablation driven mass loss. he fiux of Lyman contimmiun incident on thecloud. F=Qlid? is balanced by. the recombination rate between he Lfrout and the ilhuuinating star.," Once ionization equilibrium has been established in the ionized flow, and the system is in a steady state of photo-ablation driven mass loss, the flux of Lyman continuum incident on thecloud, $F = Q/4 \pi d^2$ is balanced by the recombination rate between the I-front and the illuminating star."
+ For a constant velocity. splicrically divergent flow. e(r)=aylryfr)? where s; is the plasma deusitv just outside (ou the radiated side) of the ionization front. aud r; is the ionization frout radius.," For a constant velocity, spherically divergent flow, $n_e(r) = n_I(r_I/r)^2$ where $n_I$ is the plasma density just outside (on the irradiated side) of the ionization front, and $r_I$ is the ionization front radius."
+" The total recombination rate in the «ωμά between the O star aud the [πο is then eiven bv where the iutegration Ην run from c, to infinity aud apm2.6«10Pens1) is the case-B recombination coefficient for lydrogen at LO! K. For a disk or circumstellar cuvironment with a mean neutral gas deusity i(H) (Ffopri)/? inside radius rj. the electron density at the L£frout is sclfreeulated by the Lviuui οὐαι flux to be ny."," The total recombination rate in the column between the O star and the I-front is then given by where the integration limits run from $r_I$ to infinity and $\alpha_B \approx 2.6 \times 10^{-13} (cm^3 s^{-1})$ is the case-B recombination coefficient for hydrogen at $10^4$ K. For a disk or circumstellar environment with a mean neutral gas density $n(H) >> (3F/\alpha_B r_I)^{1/2}$ inside radius $r_I$, the electron density at the I-front is self-regulated by the Lyman continuum flux to be $n_I$."
+ A higher electrou density would shut off the UV. decrease the nass loss. and therefore lower the deusitv.," A higher electron density would shut off the UV, decrease the mass loss, and therefore lower the density."
+ A lower electron deusity would result in a higher fiux at the I-frout aud a ereater mnass-loss rate until the deusitv at the Lfrout reached the value ny., A lower electron density would result in a higher flux at the I-front and a greater mass-loss rate until the density at the I-front reached the value $n_I$.
+ If the mean neutral eas density is lower than vy. the E-frout would fully ionize the circumstellar euvironnient which would then be optically thin to Lxiuan continuum.," If the mean neutral gas density is lower than $n_I$, the I-front would fully ionize the circumstellar environment which would then be optically thin to Lyman continuum."
+ This simple model ignores the expected increase in velocity of the photo-ablating plasma due to the density aud pressure eracieut., This simple model ignores the expected increase in velocity of the photo-ablating plasma due to the density and pressure gradient.
+ Iu a full treatment of the lvdvodvuamues. the flow velocity will increase to about 2t03 ej; as the deusity decreases by one to two orders of magnitude.," In a full treatment of the hydrodynamics, the flow velocity will increase to about 2 to 3 $\times c_{II}$ as the density decreases by one to two orders of magnitude."
+ However this acceleration will have little effect ou the observed surface brightuess of species sucli as Pao because most ofthe line eimissiou is produced by he deuse plasma at the base of the ionized flow. since ine Intensities are proportional to the cussion measure.," However this acceleration will have little effect on the observed surface brightness of species such as $\alpha$ because most ofthe line emission is produced by the dense plasma at the base of the ionized flow, since line intensities are proportional to the emission measure."
+" The emission measure is defined in the usual meamucr obe EAL=fi26)dl where n, ds the electron deusity and (f is an infinitesimal leugth element along the Ime-ot-sieht.", The emission measure is defined in the usual manner to be $EM = \int n^2_e(l)dl$ where $n_e$ is the electron density and $dl$ is an infinitesimal length element along the line-of-sight.
+ For a homisphlierical or spherical surface giviug rise oa coustaut velocity wind with au r7? deusity profile. he highest EXD occurs along a line-ofsight that just erazes the ionization frout.," For a hemispherical or spherical surface giving rise to a constant velocity wind with an $r^{-2}$ density profile, the highest EM occurs along a line-of-sight that just grazes the ionization front."
+ By iuteerating the liuc-of-sieht alone a tangent to the surface. it can be shown that EM=(s/2|peuir where pe=3.086«101?cn couverts the units of EM from cgs.," By integrating the line-of-sight along a tangent to the surface, it can be shown that $EM = (\pi/2[pc])n^2_I r_I$ where $pc = 3.086 \times 10^{18} {\rm cm}$ converts the units of EM from c.g.s."
+ to ciiSpe., to $\rm {cm^{-6} pc}$.
+" Frou ionization equilibrimin. ""ir=3O/lrageP. indepeudeut of the ionization front radius so that EAL=309/ns|pe|ojd."," From ionization equilibrium, $n^2_I r_I \approx 3Q/4\pi \alpha_B d^2$, independent of the ionization front radius, so that $EM \approx 3Q/8[pc]\alpha_B d^2$."
+ The approximate equality is used ere because we ignore he sheht acceleration of the ionized front due to the xessure eradicut., The approximate equality is used here because we ignore the slight acceleration of the ionized front due to the pressure gradient.
+ Thus. the expected surface brightucss of any emission line is. to first order. indepeudent of the cloud radius and inversely depeucdent on the square of he distance from the source of ionization.," Thus, the expected surface brightness of any emission line is, to first order, independent of the cloud radius and inversely dependent on the square of the distance from the source of ionization."
+ The dux of any recombination line of hwdrogen cau ο converted iuto au estimate of the emission measure., The flux of any recombination line of hydrogen can be converted into an estimate of the emission measure.
+ Following Spitzer(1978).. the intensity of a transition roni level τι to level an is given by where μη is the total recombination ito the upper level of the transition taking account of all reconibinatious iuto hieher levels that cascade iuto the level of interest and of the brauchiug ratio for transitions to lower levels.," Following \citet{Spit78}, the intensity of a transition from level m to level n is given by where $\alpha_{mn}$ is the total recombination into the upper level of the transition taking account of all recombinations into higher levels that cascade into the level of interest and of the branching ratio for transitions to lower levels."
+ EM is the emission measure in units of eii.Spe., EM is the emission measure in units of ${\rm cm^{-6} pc}$.
+" For full jouization. n,=nyWT) is asstuned."," For full ionization, $n_e = n_p = n(H)$ is assumed."
+ The tutensities of various hydrogen recombination lines relative to 11.) were tabulated by Peneclly (1961).., The intensities of various hydrogen recombination lines relative to $\beta$ were tabulated by \citet{Peng64}. .
+ From his tables. the (απο recombination coefficieut. for Pad is 045= 1) ," From his tables, the Case-B recombination coefficient for $\alpha$ is $\alpha_{43} = 4.10 \times 10^{-14} ({\rm cm^3 s^{-1}})$ ."
+"Thus the relatiouship between Pan fux in units of creslenJaresec7 and EM is Fp,-—2.51«10I9AL where EM is iu mits of eii be."," Thus, the relationship between $\alpha$ flux in units of $\rm {erg s^{-1} cm^{-2} arcsec^{-2}}$ and EM is $F_{P\alpha} = 2.51 \times 10^{-19} EM$ where EM is in units of ${\rm cm^{-6}pc}$ ."
+ The Pan detection lanits for our three sources are eiven in Table 2 along with the exposure times., The $\alpha$ detection limits for our three sources are given in Table 2 along with the exposure times.
+ Table, Table
+more realistic for the Suu. three of our sauple may have about solar composition and the remainius three mav be about ddex more metal rich.,"more realistic for the Sun, three of our sample may have about solar composition and the remaining three may be about dex more metal rich."
+ This means that the proportion of metal rich objects with the highest |Fo/TI]| of about |0.2 is quite high in our present sample of brown dwarts., This means that the proportion of metal rich objects with the highest [Fe/H] of about +0.2 is quite high in our present sample of brown dwarfs.
+ We have shown that € O abundances have siguificaut effects on the 5.042 spectra of brown dwirfs. and we now cxamine their effect on the 2.5 nu. spectra;," We have shown that C O abundances have significant effects on the $\mu$ m spectra of brown dwarfs, and we now examine their effect on the $\mu$ m spectra."
+" As exiuuples. we compare the predicted pau spectra of the models of UCALa and [οο series for the case of Ti,= τον. Tie= Look. and loss— LA. toygetherwiththose forthe2.5 san spectra in Figure 6.."," As examples, we compare the predicted $\mu$ m spectra of the models of UCM-a and UCM-c series for the case of $T_{\rm cr} = 1700$ K, $T_{\rm eff} = 1500$ K, and $g$ = 4.5, together with those for the $\mu$ m spectra in Figure \ref{fig:nir}."
+ The major difference between the UCALa and ο series is that IIO bands at 1.1. 1. 1.9. aud son are all stronecr iu the UCALa (curve 2 in Figure 6)) than in the UCALe series (curve 1). aud this is due to a direct effect ofthe increased oxvgen abundauce (see Figure 2)).," The major difference between the UCM-a and UCM-c series is that $_2$ O bands at 1.1, 1.4, 1.9, and $\mu$ m are all stronger in the UCM-a (curve 2 in Figure \ref{fig:nir}) ) than in the UCM-c series (curve 1), and this is due to a direct effect of the increased oxygen abundance (see Figure \ref{fig:fig2}) )."
+ Thus the IIO baud streneths depend seusitively onu oxveen abundance. and we may hope to determine oxvecu abundance from the H2O bauds.," Thus the $_2$ O band strengths depend sensitively on oxygen abundance, and we may hope to determine oxygen abundance from the $_2$ O bands."
+ Iowever. we must remember that the ΠΟ baud streneths also depeud on other parameters such as Το. Tog. and loge.oo andweshouldencounterthesa nied:f HiaPhe SEES or byotherauthorst ," However, we must remember that the $_2$ O band strengths also depend on other parameters such as $T_{\rm cr}$, $T_{\rm eff}$, and $g$, and we should encounter the same difficulty due to a degeneracy of the parameters as noted before by other authors \citep[e.g.][]{Burgasser06b, Leggett09}."
+as a metallicitv (C O abundances) indicator iu brown dwarts., This fact reconfirms the unique role of $_2$ as a metallicity (C O abundances) indicator in brown dwarfs.
+ Thanks to the spectra. we are for the first time able to demoustrate that the metallicity. more specifically € O abundances. are nmuportauf parameters to understand brown dwarf atinospheres.," Thanks to the spectra, we are for the first time able to demonstrate that the metallicity, more specifically C O abundances, are important parameters to understand brown dwarf atmospheres."
+ Until now. we have assumed. that if was sufficicut to use oue sequence of model photospheres based on a represcutative chemical composition in analyzing low resolution spectra of cool cawarts.," Until now, we have assumed that it was sufficient to use one sequence of model photospheres based on a representative chemical composition in analyzing low resolution spectra of cool dwarfs."
+ We must now aciuit that such an assunrption is duappropriate. aud we should consider abundance effects more carefully. especially of € O. in our future analysis of cool dwarfs.," We must now admit that such an assumption is inappropriate, and we should consider abundance effects more carefully, especially of C O, in our future analysis of cool dwarfs."
+ Also. we cannot use av solar composition for cool dwarfs unless this substitution can be justified by a direct analysis of the spectra of cool cdvarfs.," Also, we cannot use any solar composition for cool dwarfs unless this substitution can be justified by a direct analysis of the spectra of cool dwarfs."
+ It is true that a detailed abuudauce analysis of brown dwarf is difficult especially with low resolution spectra. but well defined molecular bands. even at low resolution. can be potential abundance indicators.," It is true that a detailed abundance analysis of brown dwarfs is difficult especially with low resolution spectra, but well defined molecular bands, even at low resolution, can be potential abundance indicators."
+ We know alreacky that CO» is a fine incicator of € ο abunudances., We know already that $_2 $ is a fine indicator of C O abundances.
+" Unfortunately, however. CO» is accessible ouly 10111 space telescopes acl. nioreover. s)octroscople observations idu the near infrared are imnostlv reelected bv the recent space iufrared missions."," Unfortunately, however, $_2$ is accessible only from space telescopes and, moreover, spectroscopic observations in the near infrared are mostly neglected by the recent space infrared missions."
+ From) the view point of the study ou cool dwairfs (and other cool stars). the iniportauce of observing he near infrared spectra (especially between 2.5 and jiu) from space cannot be emphasized too uuch.," From the view point of the study on cool dwarfs (and other cool stars), the importance of observing the near infrared spectra (especially between 2.5 and $\mu$ m) from space cannot be emphasized too much."
+ Although the spectra of brown dwirfs appear o be complicated. we are now conviuced that he spectra of brown dwirfs can basically be understood. ou the basis of the LTE mode xXhotosphlieres. but oulv if the chemical composition is properly considered.," Although the spectra of brown dwarfs appear to be complicated, we are now convinced that the spectra of brown dwarfs can basically be understood on the basis of the LTE model photospheres, but only if the chemical composition is properly considered."
+ This d$ a reasonable result for such high density photosphneres as of own dwarfs m which frequent collisious casily naintain thermal equilibrimu., This is a reasonable result for such high density photospheres as of brown dwarfs in which frequent collisions easily maintain thermal equilibrium.
+ Thus the chemica conrposifion is the most important ingredien oei the interpretation and analysis of even low --esolutiou spectra., Thus the chemical composition is the most important ingredient in the interpretation and analysis of even low resolution spectra.
+ Now. with better observe ata for brown dwarfs including those from denos yp icifeilftieteriunation cau be done iteratively for brown warts as for ordinary stars.," Now, with better observed data for brown dwarfs including those from space, analysis of the spectra and abundance determination can be done iteratively for brown dwarfs as for ordinary stars."
+" Finally, we dgnust remember that a major ifheultv in the analvsis of the spectra of brown warts is that we have no model of comparable accuracy as for ordinary stars vet."," Finally, we must remember that a major difficulty in the analysis of the spectra of brown dwarfs is that we have no model of comparable accuracy as for ordinary stars yet."
+ For this reason. even the accurate ummerical method such as outlined iu Section ?? cannot be infallible.," For this reason, even the accurate numerical method such as outlined in Section \ref{sec:fitcmp} cannot be infallible."
+ Iu fact. we have no model reproducing all the observable features correctly. aud the model found by the munerical method as well as by the eve-fitting method may prove incorrect even if they are relatively satisfactory among the models currently available.," In fact, we have no model reproducing all the observable features correctly, and the model found by the numerical method as well as by the eye-fitting method may prove incorrect even if they are relatively satisfactory among the models currently available."
+ Within this limitation. we hope that our main results on the differential effects of € ο abundances are relatively free of preseut brown chwarf model wacertaiuties.," Within this limitation, we hope that our main results on the differential effects of C O abundances are relatively free of present brown dwarf model uncertainties."
+ We thank an auonviuous referee for critical reading of the text aud for invaluable sugeestions, We thank an anonymous referee for critical reading of the text and for invaluable suggestions
+ We thank an auonviuous referee for critical reading of the text aud for invaluable sugeestionsO, We thank an anonymous referee for critical reading of the text and for invaluable suggestions
+ We thank an auonviuous referee for critical reading of the text aud for invaluable sugeestionsOO, We thank an anonymous referee for critical reading of the text and for invaluable suggestions
+transverse {ο the line of sight.,transverse to the line of sight.
+ In this case. £ should be small.," In this case, $\xi$ should be small."
+ The direction of By must be considered an unknown parameter., The direction of $\vec{B}_0$ must be considered an unknown parameter.
+ Although there is information on the form of the global Galactic magnetic field [rom Faraclay rotation measurements of pulsars and extragalactic radio sources (andandKulkarni1989:RandanclLyne1994:MinterandSpangler1996:VanEcketal2011) as well as measurements of the polarized Galactic svnchrotron emission (Becketal1996:Haverkornοἱ2004a).. (he magnetic field models are deduced from measurements on lines of sight which are kiloparsecs in leneth.," Although there is information on the form of the global Galactic magnetic field from Faraday rotation measurements of pulsars and extragalactic radio sources \citep{Rand89,Rand94,Minter96,VanEck11} as well as measurements of the polarized Galactic synchrotron emission \citep{Beck96,Haverkorn04a}, the magnetic field models are deduced from measurements on lines of sight which are kiloparsecs in length."
+ All analvses of the galactic magnetic field agree that the fIuctuating (presumably turbulent) component of the galactic magnetic [iekl is comparable to or larger than the svstematic component (e.g.BandandKulkarni1989:MinterSpaneler1996:Haverkornet2004a;:VanEcketal 2011).," All analyses of the galactic magnetic field agree that the fluctuating (presumably turbulent) component of the galactic magnetic field is comparable to or larger than the systematic component \citep[e.g.][]{Rand89,Minter96,Haverkorn04a,VanEck11}."
+. By svstematic component. we mean a vector field which is describable by a relatively simple function of Galactocentrie coordinates.," By systematic component, we mean a vector field which is describable by a relatively simple function of Galactocentric coordinates."
+" In a verv local sense. the ""large scale magnetic field is almost certainly dominated by these turbulent fIuctuations."," In a very local sense, the “large scale” magnetic field is almost certainly dominated by these turbulent fluctuations."
+ It may reflect the random orientation of the largest eddy in the solar neighborhood., It may reflect the random orientation of the largest eddy in the solar neighborhood.
+ Ii any case. Dj. and (the unit vector in the direction of By. 5. can point in anv direction in the sky.," In any case, $\vec{B}_0$, and the unit vector in the direction of $\vec{B}_0$, $\hat{b}$, can point in any direction in the sky."
+ Belore proceeding further. it is necessary (o note that there are (wo independent. aud incompatible estimates of the direction of b.," Before proceeding further, it is necessary to note that there are two independent, and incompatible estimates of the direction of $\hat{b}$."
+ Lallementetal(2005) use the difference between the direction of neutral helium flow and that of the largest concentration of neutral hvdrogen outside the to infer that 7 points in the range /=205—240°. b=—60*——38*. Gur," \cite{Lallement05} use the difference between the direction of neutral helium flow and that of the largest concentration of neutral hydrogen outside the to infer that $\hat{b}$ points in the range $l = 205^{\circ} - 240^{\circ}$, $b = -60^{\circ} - -38^{\circ}$."
+nettetal(2006). report measurements of the direction to sources of low Irecquency radio emission. which are assumed to be generated in the heliosheath?.," \cite{Gurnett06} report measurements of the direction to sources of low frequency radio emission, which are assumed to be generated in the heliosheath."
+. Thev assume that (his racio emission is generated at points on the heliopause which are perpendicular to na, They assume that this radio emission is generated at points on the heliopause which are perpendicular to $\vec{B}_0$.
+ Thev do not retrieve the vector b. but report the angle between b and the direction to the ecliptic pole.," They do not retrieve the vector $\hat{b}$, but report the angle between $\hat{b}$ and the direction to the ecliptic pole."
+ They interpret their results as being consistent will Lallementοἱal(2005)., They interpret their results as being consistent with \cite{Lallement05}.
+ A model-dependent estimate of b has also been presented by Opher, A model-dependent estimate of $\hat{b}$ has also been presented by \cite{Opher09}.
+etal(2009).. use Vovager 1 measurements of (he plasma flow direction in the heliosheath. together with an MIID model of the heliosheath. to infer that b points in the approximate direction /=10°—20°. b=28°—38°.," \cite{Opher09} use Voyager 1 measurements of the plasma flow direction in the heliosheath, together with an MHD model of the heliosheath, to infer that $\hat{b}$ points in the approximate direction $l = 10^{\circ} - 20^{\circ}$, $b = 28^{\circ} - 38^{\circ}$."
+ The estimates of Lallementetal(2005) and Opherοἱal(2009) appear to be significantly different., The estimates of \cite{Lallement05} and \cite{Opher09} appear to be significantly different.
+ It is worth emphasizing that all of (he aforementioned techniques are model-dependent in (hat they adopt physical assumptions, It is worth emphasizing that all of the aforementioned techniques are model-dependent in that they adopt physical assumptions
+variability with the source counts varying by more than a factor of two.,variability with the source counts varying by more than a factor of two.
+ FGMO03 found no evidence of variability from V710 Tau during the 50 ks 2000 oobservations., FGM03 found no evidence of variability from V710 Tau during the 50 ks 2000 observations.
+ Due to the faintness of the source. PN and MOS data were fitted simultaneously with a single temperature model: Table 9 reports the spectral parameters.," Due to the faintness of the source, PN and MOS data were fitted simultaneously with a single temperature model; Table \ref{tab:v710_all11pnmos12} reports the spectral parameters."
+ The merged PN spectrum of V710 Tau has sufficient statistics for a 2T fit., The merged PN spectrum of V710 Tau has sufficient statistics for a 2T fit.
+" The spectral parameters (N(H)-0.1920.06x10°? em7. KT,=0.30€0.053 kT»=0.93+0.06. Z=0.2 frozen) showed somewhat lower temperatures than in the 2000 observation (Τι=0.63+0.05. kT>=1.24x0.10. FGMO3)."," The spectral parameters $=0.19\pm 0.06 \times
+10^{22}$ $^{-2}$, $kT_1 = 0.30\pm 0.053$ $kT_2=0.93\pm 0.06$, $Z=0.2$ frozen) showed somewhat lower temperatures than in the 2000 observation $kT_1=0.63\pm 0.05$, $kT_2 = 1.24\pm 0.10$, FGM03)."
+ However. the SAS V6.0.0 reprocessing of the FGMO3 data flags most of the source photons as invalid (due to a row of hot pixels).," However, the SAS V6.0.0 reprocessing of the FGM03 data flags most of the source photons as invalid (due to a row of hot pixels)."
+" The reprocessed MOS? data of V710 Tau are. on the other hand. ‘clean’. and the best-fit to the 2000 data gives NCH)=0.3040.08x107 em. AT,=0.38£0.10. KT»=1.120.06 (P= 0.06). similar to the values of the current observation."," The reprocessed MOS2 data of V710 Tau are, on the other hand, `clean', and the best-fit to the 2000 data gives $=0.30\pm 0.08 \times
+10^{22}$ $^{-2}$, $kT_1 = 0.38\pm 0.10$, $kT_2=1.12\pm 0.06$ $P=0.06$ ), similar to the values of the current observation."
+" The source X-ray luminosity from the merged data is 0.6x10°""s7!.. a factor of 2 lower than the 2000 value (1.3x10 ντ)."," The source X-ray luminosity from the merged data is $0.6\times
+10^{30}$, a factor of 2 lower than the 2000 value $1.3\times
+10^{30}$ )."
+ XZ Tau is a binary CTTS with 0.3 aresee separation (Haas 1990)). associated. together with HL Tau. with a complex set of bipolar jets and Harbig Haro outflows (Mundt 1990)).," XZ Tau is a binary CTTS with 0.3 arcsec separation \citealp{hlz90}) ), associated, together with HL Tau, with a complex set of bipolar jets and Harbig Haro outflows \citealp{mbs+90}) )."
+ The spectral types are M2 and M3.5 for XZ Tau North and XZ Tau South. respectively (Hartigan&Kenyon. 2003)).," The spectral types are M2 and M3.5 for XZ Tau North and XZ Tau South, respectively \citealp{hk03}) )."
+ A photometric period of 2.6 days has been derived by Bouvieretal. (1995)., A photometric period of 2.6 days has been derived by \cite{bck+95}.
+. They interpret this period as due to rotational modulation by a hot spot. 1500 K hotter than the photosphere and covering of the stellar surface.," They interpret this period as due to rotational modulation by a hot spot, 1500 K hotter than the photosphere and covering of the stellar surface."
+ During the 2000 oobservation (FGMO3). the X-ray count rate increased by a factor of four in an approximately linear fashion over 50 ks.," During the 2000 observation (FGM03), the X-ray count rate increased by a factor of four in an approximately linear fashion over 50 ks."
+ A time-resolved spectral analysis of the X-ray emission resulted, A time-resolved spectral analysis of the X-ray emission resulted
+uagnetic mininnun of evcle 23. there are 18 N-class Hares occurred,"magnetic minimum of cycle 23, there are 18 X-class flares occurred."
+ Additionally. the anual averaged solar radio fiux Gu panel b) at frequency of 2.81 GIIz are very simular to the profile of the aunual C-class flare Munbers.," Additionally, the annual averaged solar radio flux (in panel b) at frequency of 2.84 GHz are very similar to the profile of the annual C-class flare numbers."
+" Cenerally, the solar radio cussion at 2.81 Cz is mainly associated with the non-thermal eruptive xocesses,"," Generally, the solar radio emission at 2.84 GHz is mainly associated with the non-thermal eruptive processes."
+ It shows that all of the 6 siper-flares (>V10) occurred in the decay phase of the Sclavabe cycle 23. including the the largest flare event (N28. 2003-11-03) recorded in NOAA so far.," It shows that all of the 6 super-flares $>X10$ ) occurred in the decay phase of the Schwabe cycle 23, including the the largest flare event (X28, 2003-11-03) recorded in NOAA so far."
+ These facts imply that the stronger flare events are inclined to occur in the decay phase of the Schwabe cycle., These facts imply that the stronger flare events are inclined to occur in the decay phase of the Schwabe cycle.
+ However. so far we cau not make a doubtless conclusion because we lave no enough reliable observation data of the other Sclavahbe eveles.," However, so far we can not make a doubtless conclusion because we have no enough reliable observation data of the other Schwabe cycles."
+ The above characteristics of the solar fare cistributioulil in solar Sclavabe cycles is also implied the asvuuuetric properties of the Sclavabe ceveles;, The above characteristics of the solar flare distribution in solar Schwabe cycles is also implied the asymmetric properties of the Schwabe cycles.
+ However. it is most intricate that most of the powerful flares are occurred uot in the risine-phase but iu the decay-phase.," However, it is most intricate that most of the powerful flares are occurred not in the rising-phase but in the decay-phase."
+ Tn this work. we make au assunmptiou that future behaviors of the solar activity in upcoming several decades vears can be deduced from the averaged behaviors in the past several hundred vears.," In this work, we make an assumption that future behaviors of the solar activity in upcoming several decades years can be deduced from the averaged behaviors in the past several hundred years."
+ Aud this can be accepted because the several huudred-xear is so much short related to the life-time of the Sun when it is at the main sequence., And this can be accepted because the several hundred-year is so much short related to the life-time of the Sun when it is at the main sequence.
+ We may regard the Siu as a steady-going svstem which will rui according to its averaged behaviors of its past several hundred vears and last for several hundred years again., We may regard the Sun as a steady-going system which will run according to its averaged behaviors of its past several hundred years and last for several hundred years again.
+ The time-scales of the solar 1l-vr Schlwabe cycle and 103-vr secular eveles are mmeh shorter than the dciffusiou time-scale of solar large-scale elobal magnetic field structures (about 3«107 vr. Stix. 2003). aud much arecr than the solar dvuanical time-scale (for example. he Sanin oscillations. ete).," The time-scales of the solar 11-yr Schwabe cycle and 103-yr secular cycles are much shorter than the diffusion time-scale of solar large-scale global magnetic field structures (about $3\times10^{7}$ yr, Stix, 2003), and much larger than the solar dynamical time-scale (for example, the 5-min oscillations, etc)."
+ The origin of these cvcles is à remarkable unsolved tautaliziug problem., The origin of these cycles is a remarkable unsolved tantalizing problem.
+ We need o explore some theoretical models which iav carry he energy from the interior to the surface aud release in solar afinosphere mi some periodic forms., We need to explore some theoretical models which may carry the energy from the interior to the surface and release in solar atmosphere in some periodic forms.
+ Prescutly. here are two kinds of theoretical models for solar eveles. oue is turbulent dynamo models. the other is AIIID oscillatory models.," Presently, there are two kinds of theoretical models for solar cycles, one is turbulent dynamo models, the other is MHD oscillatory models."
+ However. both of them have uuch difficulties to interpret the main features of solar eveles reliably (Tivemath. 2009).," However, both of them have much difficulties to interpret the main features of solar cycles reliably (Hiremath, 2009)."
+ As for the Ll-vr Sclawahe evele. in spite of 1iuchi of difficulties. possibly dvuame theory is the most popular hypothesis to explain the ecucrationus (Babcock. 1961: Leighton. 1961. 1969).," As for the 11-yr Schwabe cycle, in spite of much of difficulties, possibly dynamo theory is the most popular hypothesis to explain the generations (Babcock, 1961; Leighton, 1964, 1969)."
+ This model finds its orem iu the tachocline where strong shearing motions occur iu the, This model finds its origin in the tachocline where strong shearing motions occur in the
+7?7)).,\ref{sec:disrupt}) ).
+ The densitv and ionization of the ISM at the solar location are found from radiative transfer (RT) models that are constrained by observations of ISM both inside and outside of the heliosphereSFO2., The density and ionization of the ISM at the solar location are found from radiative transfer (RT) models that are constrained by observations of ISM both inside and outside of the heliosphere.
+"FS03.FS05.SF06).. While a range of tested equilibrium models reproduce the general properties of tenuous clouds. the models which provide the best fits to the local interstellar cloud (LIC) data. have n(II9)90.2I. n(e)9-0.1 ?.. 0.015 ὃν, and fractional ionizations \ 90.29. \ ~0.47.7 However somewhat lower ionizalions. n(e)e-0.077. can not vet be ruled out conclusively (FS05)."," While a range of tested equilibrium models reproduce the general properties of tenuous clouds, the models which provide the best fits to the local interstellar cloud (LIC) data, have $\sim$0.2, $\sim$ 0.1, $\sim$ 0.015 and fractional ionizations $\chi$ $\sim$ 0.29, $\chi$ $\sim$ However somewhat lower ionizations, $\sim$ 0.07, can not yet be ruled out conclusively (FS05)."
+ The Ulvsses ddata eive a LIC temperature at the heliosphere of 630025340 Ix.2004).. If observed from a distance. the warm partially ionized local ISM within ~35 pe. (IDE10P7. may appear as either warm ionized or warm neutral material (WIM. WNM).," The Ulysses data give a LIC temperature at the heliosphere of $\pm$ 340 K. If observed from a distance, the warm partially ionized local ISM within $\sim$ 35 pc, $\lesssim 10^{19} $, may appear as either warm ionized or warm neutral material (WIM, WNM)."
+ Interstellarierstellar dust. grain. abundancesabunelar in generigeneric warnrm ISM are relatively invariantinvarian {1je ionization level of the gas. and approximately of the infrared emission from dust within ~150 pe is formed in warm egas.," Interstellar dust grain abundances in generic warm ISM are relatively invariant the ionization level of the gas, and approximately of the infrared emission from dust within $\sim$ 150 pc is formed in warm gas."
+ This is shown by observations of WNM and WIM in a high-Iatitude region wilh sparse ISAL. where comparisons are made between 1001000 DDIRBE and FIRAS infrared emission. 22]-cm emission (Leiden data). and eenmission2000).," This is shown by observations of WNM and WIM in a high-latitude region with sparse ISM, where comparisons are made between 100—1000 DIRBE and FIRAS infrared emission, 21-cm emission (Leiden data), and emission."
+. Column densities for LIC-like clouds. LO 7... are not usually resolved in 22]-cem or nimneasiurenients.," Column densities for LIC-like clouds, $ < 10^{18} $ , are not usually resolved in 21-cm or measurements."
+ The LIC velocity. ως20 ii (he Local Standard of Rest2002).. is a factor in dust. composition. since high cloud velocities are associated wilh enhanced gas-phase abundances of refractory elements and erain destruction in interstellar shocks1996).," The LIC velocity, $\sim$ 15–20 in the Local Standard of Rest, is a factor in dust composition, since high cloud velocities are associated with enhanced gas-phase abundances of refractory elements and grain destruction in interstellar shocks."
+. For comparison. over of the mnmass detected in the Arecibo Millenium survey is at. LO L|.," For comparison, over of the mass detected in the Arecibo Millenium survey is at $>$ 10 ."
+ Local interstellar gas.therefore. appears to be typical dust-containing tenuous intermediate velocity partially ionized 15M.," Local interstellar gas,therefore, appears to be typical dust-containing tenuous intermediate velocity partially ionized ISM."
+ , 
+a single with a physical basis there is little hope of understanding the phenomenon.,a single with a physical basis there is little hope of understanding the phenomenon.
+ Recent observations. mainly by speckle interferometry aud by the Iipparcos satellite. have detected the binary nature of several candidates aud other candidates have been recognized on the basis of high resolution Spectra.," Recent observations, mainly by speckle interferometry and by the Hipparcos satellite, have detected the binary nature of several candidates and other candidates have been recognized on the basis of high resolution spectra."
+ We aake the hyvothesis that abuudauce zuonmalies are. at least partly. due to the effect of veiling in a coniposite spectrun aud that other still uudetected binaries are likely to be resent annue the ojects collected in Table 1.," We make the hypothesis that abundance anomalies are, at least partly, due to the effect of veiling in a composite spectrum and that other still undetected binaries are likely to be present among the objects collected in Table 1."
+ Distorted arc uncertain colours (0.9. stars with iegative colour excess and reddened bright stars) and peculiar Baliner line profiles are reasons for suspecting duplicity., Distorted and uncertain colours (e.g. stars with negative colour excess and reddened bright stars) and peculiar Balmer line profiles are reasons for suspecting duplicity.
+ The photometrically derived atmospheric parameters of close visual binaries refer to the average photometric indices of the components aud an abundance analysis based on them requires that the two compoucuts form a twin pair or have very differcut Iuuiuosities., The photometrically derived atmospheric parameters of close visual binaries refer to the average photometric indices of the components and an abundance analysis based on them requires that the two components form a twin pair or have very different luminosities.
+ Moreover. for some of these binaries the angular separation is such that a composite προςται cannot be avoided.," Moreover, for some of these binaries the angular separation is such that a composite spectrum cannot be avoided."
+Tt has been demoustrated for S-type asteroids that visible spectral slopes can be successfully used. to model the optical alterations of asteroidal surfaces due to space weatherme.,It has been demonstrated for S-type asteroids that visible spectral slopes can be successfully used to model the optical alterations of asteroidal surfaces due to space weathering.
+ Spectral slopes have been evaluated using different methods. mainly photometric data (e.g.Jedickeetal.2001) aud spectroscopy (Marchietal.20062:Paolicchict 2007).. leading basically to the sane Concerning V-types. spectroscopic data are available ouly for a restricted umber (στο) of objects. preventing auy significant characterization of the whole population.," Spectral slopes have been evaluated using different methods, mainly photometric data \citep[e.g.][]{jed04} and spectroscopy \citep{mar06a,pao07}, leading basically to the same Concerning V-types, spectroscopic data are available only for a restricted number $\sim$ 70) of objects, preventing any significant characterization of the whole population."
+ For this reason. iu the preseut work. we will make use of SDSS photometry reported in the Moving Object Catalog V3 (MOCYV23:Iveziéetal.2001:Juric2002).. which allows the aualvsis to be extended to πομιας a large umber of objects (388 and. 166 family aud nou-faily V-tvpes. respectively).," For this reason, in the present work, we will make use of SDSS photometry reported in the Moving Object Catalog V3 \citep[MOCV3;][]{ive01,jur02}, which allows the analysis to be extended to include a large number of objects (388 and 466 family and non-family V-types, respectively)."
+ V-tvpes have been selected according to Roig&CibIIluttou (2006)., V-types have been selected according to \cite{roi06}.
+". In addition. we also include 52 V-tvpes identifiedbv the spectroscopic surveys SMÁSS. .. and SINEO (Lazzaroctal.2001:Bus""S which are not included im the SDSS AIOC'V3."," In addition, we also include 52 V-types identified by the spectroscopic surveys SMASS, S3OS2, and SINEO \citep{laz04,bus02,bin04,mar05b} which are not included in the SDSS MOCV3."
+ It is iniportant to underline that the list established by Roig&Cil-IIutton(2006). consists ofcandidate V-tvpos aud that spectroscopic data are needed for a robust ideutification., It is important to underline that the list established by \cite{roi06} consists of V-types and that spectroscopic data are needed for a robust identification.
+ Nevertheless. a coniparison between candidate V-tvpes and confirmed ones displavs a good match (Roig&Cül-IIutton2006).," Nevertheless, a comparison between candidate V-types and confirmed ones displays a good match \citep{roi06}."
+. Spectral slopes have been evaluated by linear best fit of the albedo using the SDSS bauds erà. zg.," Spectral slopes have been evaluated by linear best fit of the albedo using the SDSS bands g', r',i', z'."
+ One sigma errors of SDSS albedoes have beeu used as weighting factors for the fitting procedure., One sigma errors of SDSS albedoes have been used as weighting factors for the fitting procedure.
+ As for spectroscopic data. we first evaluated the albedo ine. er. i. Z bauds through cubic spline of the spectra aud then proceeded with Imear best fit.," As for spectroscopic data, we first evaluated the albedo in g', r', i', z' bands through cubic spline of the spectra and then proceeded with linear best fit."
+ Albedoes have been normalized to the 2 baud (Roig&Cil-IInttou2006)., Albedoes have been normalized to the r' band \citep{roi06}.
+. Previous analvsis of S-tvpes aud ordinary chroudrites (OCs) showed that the space weathering on asteroids is a process that. if no other alteration is present. is expected to progressively evolve over time until eveutuallv a sort of saturation is reached.," Previous analysis of S-types and ordinary chrondrites (OCs) showed that the space weathering on asteroids is a process that, if no other alteration is present, is expected to progressively evolve over time until eventually a sort of saturation is reached."
+ Moreover. its efficicney cepeuds on the location and past orbital evolution of the Tn this work. the study of space weathering is done according to the model developed im a series of papers bv Marchietal.(2006a.b) anc Paolicchietal.(2007).," Moreover, its efficiency depends on the location and past orbital evolution of the In this work, the study of space weathering is done according to the model developed in a series of papers by \cite{mar06a,mar06b} and \cite{pao07}."
+.. This tel demonstrated that the agiug. location. aud past evolution can be condeused iuto a single parameter. lamed exposure £.," This model demonstrated that the aging, location, and past evolution can be condensed into a single parameter, named exposure $E$."
+ The exposure is proportional to the dose of radiation/1on received from the Sun aud scales as where e.c are the Keplerian clements of the body. aud £ is its average collisional age (Marchietal.2006a).," The exposure is proportional to the dose of radiation/ion received from the Sun and scales as where $a, e$ are the Keplerian elements of the body, and $T$ is its average collisional age \citep{mar06a}."
+. For MBAs we used proper elemieuts eiven their stabilitv on a longer timescale., For MBAs we used proper elements given their stability on a longer timescale.
+ Ages are derived according to the collisional evolution of mainbelt asteroids (Bottkeetal.2005:Marchietal.2006a).," Ages are derived according to the collisional evolution of main belt asteroids \citep{bot05,mar06a}."
+. Collisional Lifetimes depeud. upon the diameters. which have been estimated usus the measured average ecometric albedo (0.3) of a siuuple of V-tvpoes (Tedescoetal.2002:Benneral.2002:Delbóet2003. 2006).," Collisional lifetimes depend upon the diameters, which have been estimated using the measured average geometric albedo (0.3) of a sample of V-types \citep{ted02,ben02,del03,del06}."
+. The collisional age of V-tvpes spaus from about 0.5 to 2.7 Ca. These extremes correspond to the simallest aud largest bodies of the sample. which have estimated diameters of about D km aud 10 kin. respectively.," The collisional age of V-types spans from about 0.5 to 2.7 Ga. These extremes correspond to the smallest and largest bodies of the sample, which have estimated diameters of about 1 km and 10 km, respectively."
+ Iu what follows. it is asstumed that the formation of Vestas Bunilv occurred sometime before the oldest age in the sample. namely ~2.7 Ca. so that asteroid’s age are set by subsequent collisional evolution.," In what follows, it is assumed that the formation of Vesta's family occurred sometime before the oldest age in the sample, namely $\sim2.7$ Ga, so that asteroid's age are set by subsequent collisional evolution."
+ Although Vesta’s family age is not well coustrained. a recent work of NesvoruYetal. found evidence for a mach earlier origin. possibly dating 3.5-3.8 Ca ago. than previously assumed (Alarzarietal.1996).," Although Vesta's family age is not well constrained, a recent work of \cite{nes08} found evidence for a much earlier origin, possibly dating 3.5-3.8 Ga ago, than previously assumed \citep{mar96}."
+. For the NEOs. eiven their short evolution timescale in near-Earth space (~5 Muy). the major contribution to thei space weathering alteration comes from their past evolution iuto the main belt.," For the NEOs, given their short evolution timescale in near-Earth space $\sim5$ Myr), the major contribution to their space weathering alteration comes from their past evolution into the main belt."
+ Therefore. for cach NEO we set," Therefore, for each NEO we set"
+using apertures (see Table 1)) centred on the peak emission.,using apertures (see Table \ref{table}) ) centred on the peak emission.
+" Poglitsch,Waelkens&Geis(2010) detail the flux calibration of PACS data and estimate the calibration uncertainties in the measured flux densities to be and for 100 and 160pm,, respectively."," \cite{pog10}
+ detail the flux calibration of PACS data and estimate the calibration uncertainties in the measured flux densities to be and for 100 and 160, respectively."
+ The dominant flux calibration uncertainties have been combined in quadrature with statistical uncertainties from the rms levels in the maps., The dominant flux calibration uncertainties have been combined in quadrature with statistical uncertainties from the rms levels in the maps.
+ These combined uncertainties are applied to the fluxes in Table 1.., These combined uncertainties are applied to the fluxes in Table \ref{table}.
+ The flux densities reported in Table 1 are plotted on spectral energy distributions in refSEDs.., The flux densities reported in Table \ref{table} are plotted on spectral energy distributions in \\ref{SEDs}.
+" The disc components of β Leo and 6 UMa are well fit by a simple blackbody in the absence of submillimetre detections, but 7 Corvi requires a two component fit to its disc emission: a warmer blackbody and a modified blackbody for the cold component to fit the submillimetre flux densities."," The disc components of $\beta$ Leo and $\beta$ UMa are well fit by a simple blackbody in the absence of submillimetre detections, but $\eta$ Corvi requires a two component fit to its disc emission: a warmer blackbody and a modified blackbody for the cold component to fit the submillimetre flux densities."
+" The temperature, radius (Κι) and fractional luminosity of these fits are reported in Table 1.."," The temperature, radius $R_{dust}$ ) and fractional luminosity of these fits are reported in Table \ref{table}."
+ Fitting of 2D Gaussians to each source at 100 yyields FWHM values (see Table 1)) larger than the nominal PACS PSF of 6777., Fitting of 2D Gaussians to each source at 100 yields FWHM values (see Table \ref{table}) ) larger than the nominal PACS PSF of 7.
+ Analysis of an observation of Vesta yields a PSF of 6766 x 6799., Analysis of an observation of Vesta yields a PSF of 6 $\times$ 9.
+ Vesta is a cool blackbody for which the response within the 100 ffilter should be very similar to our dust discs., Vesta is a cool blackbody for which the response within the 100 filter should be very similar to our dust discs.
+" It has a temperature measured in the submillimetre at 130-160 K (Chamberlain,Lovell&Sykes2007),, slightly warmer than our two A-star dist discs at ~110 K. Taking into account the"," It has a temperature measured in the submillimetre at 130-160 K \citep{chamber07}, slightly warmer than our two A-star dist discs at $\sim$ 110 K. Taking into account the"
+the three Drimmel&Spergel(2001) models aud the three Spergeletal.(1996) models.,the three \citet{DS01} models and the three \citet{SMB96} models.
+ For the AO. SO. and axisvanmetric analvtic magnetic fields. ihe Drimmel&Spergel(2001) models produce slightly higher normalized polarizations than the Spergeletal.(1996) models.," For the A0, S0, and axisymmetric analytic magnetic fields, the \citet{DS01} models produce slightly higher normalized polarizations than the \citet{SMB96} models."
+ DEMO magnetic fields show the opposite elfect. with the Drimmel&Spergel(2001) models producing shehtly lower normalized polarizations (han Spereeletal.(1996). based models.," DEHO magnetic fields show the opposite effect, with the \citet{DS01} models producing slightly lower normalized polarizations than \citet{SMB96} based models."
+ All of these simulations are optically thin. and so wavelength independent.," All of these simulations are optically thin, and so wavelength independent."
+ However. wavelength is known to affect the polarizing elliciency of dust grains 1975).," However, wavelength is known to affect the polarizing efficiency of dust grains \citep{SMF75}."
+. Under the assumption that there is a unilorm wavelength dependence of polarizing efficiency along the lines of sight. scaling between. wavelengths only requires use of a multiplicative [actor between the predicted. normalized degree of polarization distribution and the actual polarization distribution.," Under the assumption that there is a uniform wavelength dependence of polarizing efficiency along the lines of sight, scaling between wavelengths only requires use of a multiplicative factor between the predicted, normalized degree of polarization distribution and the actual polarization distribution."
+ These effects will not change the observed PA for a given star since. under these assumptions. (he polarizing efficiency does not alfect the PA calculation.," These effects will not change the observed PA for a given star since, under these assumptions, the polarizing efficiency does not affect the PA calculation."
+ In Fig. 2..," In Fig. \ref{CDFs},"
+ the CDFs of the P values in each model are shown. normalized by the star with the highest predicted. polarization.," the CDFs of the P values in each model are shown, normalized by the star with the highest predicted polarization."
+ Each model was normalized independently., Each model was normalized independently.
+ The star with the highest predicted polarization value is (vpically 20-40 limes the average modeled. polarization., The star with the highest predicted polarization value is typically 20-40 times the average modeled polarization.
+ In of the models. the star with the highest polarization was a sinnlated G2 supergiant al 13.31 kpe. and in of the models it was a simulated FS supereiant al 13.51 kpc (all of the models used the same stellar population samples).," In of the models, the star with the highest polarization was a simulated G2 supergiant at 13.31 kpc, and in of the models it was a simulated F8 supergiant at 13.51 kpc (all of the models used the same stellar population samples)."
+ These were some of the most distant stars returned by the Robinetal.(2003). stellar population model., These were some of the most distant stars returned by the \citet{R03} stellar population model.
+ The long distances between (hese stars and observer generated larger degrees of starlight polarization., The long distances between these stars and observer generated larger degrees of starlight polarization.
+ These bright stars are exceecinely rare. but may be detected in maenitude-limited polarization survevs covering large Ilractions of the sky.," These bright stars are exceedingly rare, but may be detected in magnitude-limited polarization surveys covering large fractions of the sky."
+ Since these high polarization stars make up such a small fraction of the total population. the CDFs shown," Since these high polarization stars make up such a small fraction of the total population, the CDFs shown"
+weak turbulence κkk.,weak turbulence $\propto k_{\perp}^{-2}$.
+" Tf however the Ay width of the forcing is limited. but we cau achieve arbitrarily high resolution in the &Á, direction. the iteraction between Alfvénn modes will eventually become strong euoueh to satisty critical balance and establish a stroug turbulence Spectruni."," If however the $k_\|$ width of the forcing is limited, but we can achieve arbitrarily high resolution in the $k_\perp$ direction, the interaction between Alfvénn modes will eventually become strong enough to satisfy critical balance and establish a strong turbulence spectrum."
+ This is partly supported by the following derivation., This is partly supported by the following derivation.
+ It can be proved that turbulent fluctuations deseribed by system (5)) satisfy the exact relation: whereóz*—z'(x|ri)z'(x).and nt—ózvafry is the longitudinal compoucut of 627.," It can be proved that turbulent fluctuations described by system \ref{restricted}) ) satisfy the exact relation: where $\delta {\tilde {\bf z}}^{\pm}= {\tilde {\bf z}}^{\pm}({\bf
+ x}+{\bf r}_{\perp})-{\tilde {\bf z}}^{\pm}({\bf x})$, and $\delta {\tilde z}^{\pm}_{l}=\delta {\tilde {\bf z}}^{\pm}\cdot
+{\bf r}_{\perp}/r_{\perp}$ is the longitudinal component of $\delta {\tilde {\bf z}}^{\pm}$."
+ Averaging is taken over the statistical cusemble. or over time and spatial positionx.," Averaging is taken over the statistical ensemble, or over time and spatial position."
+ In this formmla. €* are the coustant rates of i! and i. enerev dissipation.," In this formula, $\epsilon^{\pm}$ are the constant rates of ${\tilde z}^{+}$ and ${\tilde z}^{-}$ energy dissipation."
+ Iu the isotropic case. that is. without the euide field. the relation analogous to (6)) was derived by Politano&Pouquet(1998).," In the isotropic case, that is, without the guide field, the relation analogous to \ref{pp-restricted}) ) was derived by \citet{politano-pouquet}."
+. We now prove the following oogeneral inequality: The first step is to use the Scelwartz inequality. (82S(datyyx(a2;-D the second step is to note that (irP<ο which completes the proof.," We now prove the following general inequality: The first step is to use the Schwartz inequality, $\langle \delta {\tilde z}^{\pm}_{l}(\delta {\tilde {\bf
+ z}}^{\mp})^2 \rangle^2\leq \langle (\delta {\tilde z}_{l}^{\pm})^2 \rangle
+\langle (\delta {\tilde {\bf z}}^{\mp})^4\rangle $; the second step is to note that $(\delta {\tilde z}^{\pm}_{l})^2\leq (\delta {\tilde {\bf
+ z}}^{\pm})^2$, which completes the proof."
+ Both sides iu the expression (7)) have finite limits as viscosity and resistivity eo to zero., Both sides in the expression \ref{mhd-inequality-restricted}) ) have finite limits as viscosity and resistivity go to zero.
+ In the inertial iuterval. the correlation functions on the right-hand side of this expression have a power-law behavior. that is. ((6z*)?3xry&2 ouk(MjbXrpE (we assunie statistical sviunietrv between z! and2 }.," In the inertial interval, the correlation functions on the right-hand side of this expression have a power-law behavior, that is, $\langle (\delta {\tilde {\bf z}}^{\pm})^2\rangle
+\propto r_{\perp}^{{\zeta}_2}$ and $\langle (\delta {\tilde {\bf
+ z}}^{\pm})^4\rangle \propto r_{\perp}^{{\zeta}_4}$ (we assume statistical symmetry between ${\tilde {\bf z}}^+$ and ${\tilde {\bf z}}^-$ )."
+ Since the left haud side of (7)) scales as x»rj. uud the inequality/. should hold for. arbitrarily:. πια ki. we obtain the exact result This inequality is useful for evaluation of these exponents from nunerieal simulations or experiments since theratio of these exponents is well nieasured by the nethod of extended selfsimilarity (Benzietal.1993)..," Since the left hand side of \ref{mhd-inequality-restricted}) ) scales as $\propto r_{\perp}^2$, and the inequality should hold for arbitrarily small $r_{\perp}$, we obtain the exact result This inequality is useful for evaluation of these exponents from numerical simulations or experiments since the of these exponents is well measured by the method of extended self-similarity \citep{benzi}."
+" The inequality (8)) then provides a boundary ou the urbulence cherey spectrum that is related to the secoud-order scaling exponent as ΓΕ}xkt “Ta our case. he scaling exponent ¢, is usually close to 265 within sunall intermittency corrections. which can be checked mnercallv (e.g..Mülleretal. 2003).."," The inequality \ref{bound-mhd}) ) then provides a boundary on the turbulence energy spectrum that is related to the second-order scaling exponent as $E(k_{\perp})\propto k_{\perp}^{-1-\zeta_2}.$ In our case, the scaling exponent ${\zeta}_4$ is usually close to $2{\zeta}_2$ within small intermittency corrections, which can be checked numerically \citep[e.g.,][]{muller-biskamp-grappin}. ."
+ Inequality (83) then nuplies that ου<2. aud. therefore. the field-perpendicular cuerey spectim cannot be esseutiallv steeper thanhowever. ΓΕ)xkyj im the lit kj»x.," Inequality \ref{bound-mhd}) ) then implies that $3{\zeta}_2\leq 2$, and, therefore, the field-perpendicular energy spectrum cannot be essentially steeper than $E(k_{\perp})\propto k^{-5/3}_{\perp}$ in the limit $k_\perp \to \infty$."
+ Note. that in our muuerical &udiues the spectral exponent 5/3 is not distinguished in auv wav: rather. the field-perpeudicular energy spectrum of strong ΑΠΤΟ turbulence is flatter and closer το. 3/2.," Note, however, that in our numerical findings the spectral exponent $-5/3$ is not distinguished in any way; rather, the field-perpendicular energy spectrum of strong MHD turbulence is flatter and closer to $-3/2$."
+ This is consistent with receut results of Müller&Carap-piu(2005):Masonotal.(2007) aud also with hieh-resolutiou sinimlatious of isotropic ΑΠΟ turbulence bv IHaugenetal.(2003):Miniuni&Pouguet(2007).," This is consistent with recent results of \citet{muller,mason2} and also with high-resolution simulations of isotropic MHD turbulence by \citet{haugen,mininni}. ."
+. Astrophysical observatious of the soluw wind aud of the interstellar ποπ reveal the presence of MIID turbulence. and fud support for both 5/3 and 3/2 spectral exponeuts (e.9..Goldsteinetal.1995:etal.2006:Sunitnova 2006)..," Astrophysical observations of the solar wind and of the interstellar medium reveal the presence of MHD turbulence, and find support for both $-5/3$ and $-3/2$ spectral exponents \citep[e.g.,][]{goldstein,goldstein-roberts,bale,borovsky,podesta,smirnova}."
+ However. statistics of sucli data are often not good enough to distinguish between 5/37 aud 3/27 with confidence.," However, statistics of such data are often not good enough to distinguish between $-5/3$ ” and $-3/2$ ” with confidence."
+ Ou the nunierical side. simulations of ΑΠΟ turbulence iu the framework of reduced MIID were performed iu many works (e...Dauitruketal.2003:Gomezct2005: 2007): however. either the simulation domain Was not anisotropic to ensure the critical balance condition (3)). or the driviug force was not spatially ioniogeneonus. for example. applied at the boundary of he domain.," On the numerical side, simulations of MHD turbulence in the framework of reduced MHD were performed in many works \citep[e.g.,][]{dmitruk,gomez,rapazzo}; ; however, either the simulation domain was not anisotropic to ensure the critical balance condition \ref{crit}) ), or the driving force was not spatially homogeneous, for example, applied at the boundary of the domain."
+ Our results sugeest that the interpretation of observational aud uuinierical results may be obscured if hie Ay aud Ay structure of the spectrum is either uot well measured or not well controlled. in which case it is rard to deduce whether the field-parallel dyaizuuics have con captured and whether the universal regine of MITD urbulence has been established.," Our results suggest that the interpretation of observational and numerical results may be obscured if the $k_\|$ and $k_{\perp}$ structure of the spectrum is either not well measured or not well controlled, in which case it is hard to deduce whether the field-parallel dynamics have been captured and whether the universal regime of MHD turbulence has been established."
+ This work was supported by the US Departincnt of Enereyv under eraut DE-FEGO2-07ER51932. aud bv the NSF Center for Maguetie. Sel£-Oreuizatiou πι Laboratoryand Astrophysical Plasmas at theUniversity ofWiscousin-\acdisou.," This work was supported by the US Department of Energy under grant DE-FG02-07ER54932, and by the NSF Center for Magnetic Self-Organization in Laboratoryand Astrophysical Plasmas at theUniversity ofWisconsin-Madison."
+ IHüeh-perforuiunicc- resources were provided by the Texas Advanced Computing Center (TACC) at the University of Texas at Austin under the NSF-Teragrid Project, High-performance-computing resources were provided by the Texas Advanced Computing Center (TACC) at the University of Texas at Austin under the NSF-Teragrid Project TG-PHY070027T.
+ URL:http://www.tacc.utexas.edu, URL:http://www.tacc.utexas.edu
+use of older technology and a single filter led to various biases.,use of older technology and a single filter led to various biases.
+ These old catalogs suffered as much from being black and white as they did from being eye-selected., These old catalogs suffered as much from being black and white as they did from being eye-selected.
+" Even more disturbing, measures of completeness and contamination in the Abell catalog disagree by factors of a few."," Even more disturbing, measures of completeness and contamination in the Abell catalog disagree by factors of a few."
+" Unfortunately, some of these problems will plague any optically selected cluster sample, but the use of color information, objective selection criteria and a strong statistical understanding of the catalog can mitigate their effects."," Unfortunately, some of these problems will plague any optically selected cluster sample, but the use of color information, objective selection criteria and a strong statistical understanding of the catalog can mitigate their effects."
+ Only in the past twenty years has it become possible to utilize the objectivity of computational algorithms in the search for galaxy clusters., Only in the past twenty years has it become possible to utilize the objectivity of computational algorithms in the search for galaxy clusters.
+" These more modern studies required that plates be digitized, so that the data are in machine readable form."," These more modern studies required that plates be digitized, so that the data are in machine readable form."
+ The hybrid technology of digitized plate surveys blossomed into a cottage industry., The hybrid technology of digitized plate surveys blossomed into a cottage industry.
+" The first objective catalog produced was the Edinburgh/Durham Cluster Catalog ?),, which covered 0.5 sr (~1,600 square degrees) (EDCC,around the South Galactic Pole (SGP)."," The first objective catalog produced was the Edinburgh/Durham Cluster Catalog , which covered 0.5 sr $\sim 1,600$ square degrees) around the South Galactic Pole (SGP)."
+" Later, the APM cluster catalog was created by applying Abell-like criteria to select overdensities(?) from the galaxy catalogs."," Later, the APM cluster catalog was created by applying Abell-like criteria to select overdensities from the galaxy catalogs."
+" The largest, most recent, and the last of the photo-digital cluster survey is the Northern Sky Optical Survey???7)."," The largest, most recent, and the last of the photo-digital cluster survey is the Northern Sky Optical Survey."
+" This survey relies on galaxy catalogs created from scans of the second generation Palomar Sky Survey plates, input to an adaptive kernel galaxy density mapping routine."," This survey relies on galaxy catalogs created from scans of the second generation Palomar Sky Survey plates, input to an adaptive kernel galaxy density mapping routine."
+" The final catalog covers 11,733 square degrees, with nearly 16,000 candidate clusters, extendingto z~ 0.3."," The final catalog covers 11,733 square degrees, with nearly 16,000 candidate clusters, extendingto $z\sim0.3$ ."
+ A supplemental catalog up to z~0.5 was generated by using Voronoi Tessellation and Adaptive Kernel maps., A supplemental catalog up to $z\sim0.5$ was generated by using Voronoi Tessellation and Adaptive Kernel maps.
+" With the advent of CCDs, fully digital imaging in astronomy became reality."," With the advent of CCDs, fully digital imaging in astronomy became a reality."
+" These detectors provided an order-of-magnitudea increase in sensitivity, linear response to light, small pixel size, stability, and much easier calibration."," These detectors provided an order-of-magnitude increase in sensitivity, linear response to light, small pixel size, stability, and much easier calibration."
+" The main drawback relative to photographic plates was (and remains) their small physical size, which permits only a small area (of order to be imaged by a larger 4096? pixel detector."," The main drawback relative to photographic plates was (and remains) their small physical size, which permits only a small area (of order $15'$ ) to be imaged by a larger $4096^2$ pixel detector."
+" Realizing15’) the vast scientific potential of such a survey, an international collaboration embarked on the Sloan Digital Sky Survey (SDSS, which included construction of a specialized sdss2.5 org).meter telescope, a camera with a mosaic of 30 CCDs, a novel observing strategy, and automated pipelines for survey operations and data processing."," Realizing the vast scientific potential of such a survey, an international collaboration embarked on the Sloan Digital Sky Survey (SDSS, ), which included construction of a specialized 2.5 meter telescope, a camera with a mosaic of 30 CCDs, a novel observing strategy, and automated pipelines for survey operations and data processing."
+" Main survey operations were completed in the fall of 2005, with over 8,000 square degrees of the northern sky image in five filters to a depth of r’~22.2 with calibration accurate to ~1—2%, as well as spectroscopy of nearly one million objects."," Main survey operations were completed in the fall of 2005, with over 8,000 square degrees of the northern sky image in five filters to a depth of $r'\sim22.2$ with calibration accurate to $\sim1-2\%$, as well as spectroscopy of nearly one million objects."
+" With such a rich dataset, many groups both internal and external to the SDSS collaboration have generated a variety of cluster catalogs, from both the photometric and the spectroscopic catalogs, using techniques including: Each method generates a different catalog, and early attempts to compare them have shown not only that they are quite distinct, but also that comparison of two photometrically-derived cluster catalogs, even from the same galaxy catalog, is not straightforward(?)."," With such a rich dataset, many groups both internal and external to the SDSS collaboration have generated a variety of cluster catalogs, from both the photometric and the spectroscopic catalogs, using techniques including: Each method generates a different catalog, and early attempts to compare them have shown not only that they are quite distinct, but also that comparison of two photometrically-derived cluster catalogs, even from the same galaxy catalog, is not straightforward."
+". In addition to the SDSS, smaller areas, but to much higher redshift, have been covered by numerous deep CCD imaging surveys."," In addition to the SDSS, smaller areas, but to much higher redshift, have been covered by numerous deep CCD imaging surveys."
+" Notable examples include the Palomar Distant Cluster Survey?),, the ESO Imaging Survey?),, and many(PDCS, others."," Notable examples include the Palomar Distant Cluster Survey, the ESO Imaging Survey, and many others."
+" None of these surveys provide the angular coverage necessary for large-scale structure and precision cosmology studies, and have been specifically designed to find rich clusters at high redshift."," None of these surveys provide the angular coverage necessary for large-scale structure and precision cosmology studies, and have been specifically designed to find rich clusters at high redshift."
+" The largest such survey to date is the Red Sequence Cluster Survey?),, based on moderately deep two-band imaging using the CFH12K mosaic camera on the CFHT 3.6m telescope, covers ~100 square degrees."," The largest such survey to date is the Red Sequence Cluster Survey, based on moderately deep two-band imaging using the CFH12K mosaic camera on the CFHT 3.6m telescope, covers $\sim100$ square degrees."
+ This area coverage is comparable to X-ray surveys designed to detect clusters at z~1(?).., This area coverage is comparable to X-ray surveys designed to detect clusters at $z\sim1$.
+ Any cluster survey must make many different mathematical and methodological choices., Any cluster survey must make many different mathematical and methodological choices.
+" Regardless of the data set and algorithm used, a few simple rules should be followed to produce a catalog that is useful for statistical studies of galaxy populations and for cosmological tests: One of the most popular and commonly used methods today is the Voronoi Tesselation??7?)."," Regardless of the data set and algorithm used, a few simple rules should be followed to produce a catalog that is useful for statistical studies of galaxy populations and for cosmological tests: One of the most popular and commonly used methods today is the Voronoi Tesselation."
+ Our implementation of this technique is described in detail in §??.., Our implementation of this technique is described in detail in \ref{algorithm}.
+" Briefly, it subdivides a spatial distribution into a unique set of polygonal cells, one for each object, with the cell size inversely proportional to the local density."," Briefly, it subdivides a spatial distribution into a unique set of polygonal cells, one for each object, with the cell size inversely proportional to the local density."
+" One then defines a galaxy cluster as a high density region, composed of small adjacent cells."," One then defines a galaxy cluster as a high density region, composed of small adjacent cells."
+" Voronoi Tesselation satisfies the above criteria for generating statistical, objective, cluster samples."," Voronoi Tesselation satisfies the above criteria for generating statistical, objective, cluster samples."
+ It requires no a, It requires a
+magnetic fields (Kataoka&Stawarz2005;Crostonetal.2005).,"magnetic fields \citep{kat05,cros05}."
+". We calculate the expected Compton fluxes from the lobes on the assumption of equipartition using the code of Hardeastle,(1998) and with the same assumptions as Crostonetal.(2005) (lobes are treated as uniform cylinders, μια=10. 0=2, and a break in the electron spectrum is applied to match the radio data), taking normalizing flux densities for the lobe regions, excluding the NHS and SHS, trom unpublished archival low-frequency VLA (1.4 GHz, observation ALO146) and GMRT (610 and 244 MHz, observation number 2070) data."," We calculate the expected inverse-Compton fluxes from the lobes on the assumption of equipartition using the code of \cite{mjh98} and with the same assumptions as \cite{cros05} (lobes are treated as uniform cylinders, $\gamma_{\rm min} = 10$, $\delta=2$, and a break in the electron spectrum is applied to match the radio data), taking normalizing flux densities for the lobe regions, excluding the NHS and SHS, from unpublished archival low-frequency VLA (1.4 GHz, observation AL0146) and GMRT (610 and 244 MHz, observation number 2070) data."
+" 3C 33 1s unusual in that the flux density even at low frequencies is dominated by the bright hot spots, not the lobes."," 3C 33 is unusual in that the flux density even at low frequencies is dominated by the bright hot spots, not the lobes."
+" This calculation predicts flux densities of 0.7 nJy for cach lobe at equipartition, α factor 1.9 and 2.6 below the observed values."," This calculation predicts flux densities of 0.7 nJy for each lobe at equipartition, a factor 1.9 and 2.6 below the observed values."
+" As Crostonetal.(2005) find that this factor is typically —2 for the sources they study, 3335 lobe X-ray emission gives results consistent with those seen in other FR Ls."," As \cite{cros05} find that this factor is typically $\sim 2$ for the sources they study, 33's lobe X-ray emission gives results consistent with those seen in other FR IIs."
+ The X-ray morphology does not precisely match the radio morphology. although the low X-ray surface brightness and heavy smoothing of the data prevent a definitive statement.," The X-ray morphology does not precisely match the radio morphology, although the low X-ray surface brightness and heavy smoothing of the data prevent a definitive statement."
+ Thus it is at least plausible that this extended X-ray emission is due to inverse-Compton emission., Thus it is at least plausible that this extended X-ray emission is due to inverse-Compton emission.
+ The other possibility is that this X-ray emission is from an extended hot gas corona., The other possibility is that this X-ray emission is from an extended hot gas corona.
+" We consider two models for the distribution of the gas, a uniform density sphere and a shell surrounding the radio lobes."," We consider two models for the distribution of the gas, a uniform density sphere and a shell surrounding the radio lobes."
+ Modeling the emission region as a uniform-density sphere of radius I’ (68.4 kpe) with the best-fit temperature of the thermal model. we find a gas density of 7x  * and a total mass of 2x 10! M...," Modeling the emission region as a uniform-density sphere of radius $'$ (68.4 kpc) with the best-fit temperature of the thermal model, we find a gas density of $\times$ $^{-4}$ $^{-3}$ and a total mass of $\times$ $^{10}$ $M_\odot$."
+ A more realistic density profile (i.c. a beta-model) would only change these numbers by a factor of order unity., A more realistic density profile (i.e. a beta-model) would only change these numbers by a factor of order unity.
+" This is a relatively small amount of gas, probably less than the mass of stars in the host galaxy."," This is a relatively small amount of gas, probably less than the mass of stars in the host galaxy."
+" The optical luminosity of the host galaxy (3 ,,2-20.7) implies a stellar mass of ~9x 10!"" M. assuming a mass to light ratio of 6 (Binney&Tremaine1987).", The optical luminosity of the host galaxy $M_B$ =-20.7) implies a stellar mass of $\sim$ $\times$ $^{10}$ $M_\odot$ assuming a mass to light ratio of 6 \citep{bin87}.
+". The total thermal energy and pressure of this gas is ~1.3 10°"" ergs and ~3.2x ? dyn respectively."," The total thermal energy and pressure of this gas is $\sim$ $\times$ $^{59}$ ergs and $\sim$ $\times$ $^{-12}$ dyn $^{-2}$, respectively."
+ The thermal energy of the gas is not particularly large compared with the mechanical energy of many powerful radio galaxies such as 3C 33 (Birzanetal.2004;Kraft2006).," The thermal energy of the gas is not particularly large compared with the mechanical energy of many powerful radio galaxies such as 3C 33 \citep{bir04,kra06}."
+". Any atmosphere of this mass, or less. is likely to have been shock heated by the radio outflow, and so would no longer be bound to the relatively shallow gravitational potential of the host galaxy of 3C 22OD."," Any atmosphere of this mass, or less, is likely to have been shock heated by the radio outflow, and so would no longer be bound to the relatively shallow gravitational potential of the host galaxy of 3C 33."
+" Alternatively, the X-ray emission coincident with the lobes could be from a shell of gas that has been swept up and compressed by the inflation of the radio souree."," Alternatively, the X-ray emission coincident with the lobes could be from a shell of gas that has been swept up and compressed by the inflation of the radio source."
+" Modeling the northern emission region as a spherical shell with a radiusof 62 Κρο and thickness of 10 kpe, with the best-fitting temperature of the thermal model. we find a proton density in this shell of 6.0x10. . anda total mass of 7x10 M. ."," Modeling the northern emission region as a spherical shell with a radiusof 62 kpc and thickness of 10 kpc, with the best-fitting temperature of the thermal model, we find a proton density in this shell of $6.0 \times 10^{-3}$ $^{-3}$, and a total mass of $7
+\times 10^{8}$ $M_\odot$ ."
+" Fora thinner, 1-kpe thick shell. the mass is 2xLO* (similar"," For a thinner, 1-kpc thick shell, the mass is $2 \times 10^{8}$ (similar"
+"bilinear vyrj, relevant lor Mj transforms as We assume a scalar singleto \ (transformingo as X—wy for class Bs> and 4—«7 for class B,,6 which leads (ο the folowing Z4 invariant Yukawa Lagrangians lor classes D; and D; where o=imo"", ",bilinear $\nu_{R_j} \nu_{R_k}$ relevant for $M_R$ transforms as We assume a scalar singlet $\chi$ transforming as $\chi \rightarrow \omega \chi$ for class $B_5$ and $\chi \rightarrow \omega^2 \chi$ for class $B_6$ which leads to the following $Z_3$ invariant Yukawa Lagrangians for classes $B_5$ and $B_6$: where $\tilde{\phi} = i \tau_2 \phi^*$.
+Next. we show how a large effective neutrino mass can arise in such a model.," Next, we show how a large effective neutrino mass can arise in such a model."
+ We note that p contains (wo (vpes of mass terms viz., We note that $M_R$ contains two types of mass terms viz.
+" 1) Bare mass term which does nol need a scalar singlet ancl is invariant by itself,", 1) Bare mass term which does not need a scalar singlet and is invariant by itself.
+ 2) Terms arising from Yukawa couplings to Y., 2) Terms arising from Yukawa couplings to $\chi$.
+ The scale of latter is restricted by the scale of Z4 breaking while there is no such restriction on the bare mass term which can have a higher mass scale., The scale of latter is restricted by the scale of $Z_3$ breaking while there is no such restriction on the bare mass term which can have a higher mass scale.
+" It can be seen [rom equ.(3) that the ee and ji entries of M, have contributions to their numerators from ee and µτ entries of My which arise [rom (he bare mass term.", It can be seen from eqn.(3) that the $ee$ and $\mu \tau$ entries of $M_\nu$ have contributions to their numerators from $ee$ and $\mu \tau$ entries of $M_R$ which arise from the bare mass term.
+" We assume (he mass eigenvalues of Mp (o have same order of magnitude which leads to a large value of ee and jiz entries of M, while the other elements of Af, are suppressed. thus. leading to à large value of."," We assume the mass eigenvalues of $M_D$ to have same order of magnitude which leads to a large value of $ee$ and $\mu \tau$ entries of $M_\nu$ while the other elements of $M_\nu$ are suppressed, thus, leading to a large value of $M_{ee}$."
+" Since these textures are realized al (le seesaw scale. the Renormalization Group (IG) evolution of (he parameters of AL, [rom the seesaw scale to the electroweak scale needs to be taken into account."," Since these textures are realized at the seesaw scale, the Renormalization Group (RG) evolution of the parameters of $M_\nu$ from the seesaw scale to the electroweak scale needs to be taken into account."
+ It is well known that the RG effects are most prominent for the quasiclegenerate mass spectrum which is precisely the case here due to the assumption of large A., It is well known that the RG effects are most prominent for the quasidegenerate mass spectrum which is precisely the case here due to the assumption of large $M_{ee}$ .
+" However. it is also known that zero minors in AM,. al a given energy scale. remain zero αἱ any other energv scale at the one oop level [7]."," However, it is also known that zero minors in $M_\nu$ , at a given energy scale, remain zero at any other energy scale at the one loop level \cite{7}."
+". This is because the matrices at any (wo energy scales sry and fo are related by Af.(qui)=FM,. where f is diagonal. positive and non singular."," This is because the matrices at any two energy scales $\mu_1$ and $\mu_2$ are related by $M_\nu(\mu_1)=IM_{\nu}(\mu_2)I$, where $I$ is diagonal, positive and non singular."
+" The operation of diagonal natrices from left and right on AZ, does not alter the zero minors of M, leading to zero minors in M, at any other scale.", The operation of diagonal matrices from left and right on $M_\nu$ does not alter the zero minors of $M_\nu$ leading to zero minors in $M_\nu$ at any other scale.
+ We reconstruct the neutrino mass matrix in the flavor basis assuming neutrinos to be \Majorana particles., We reconstruct the neutrino mass matrix in the flavor basis assuming neutrinos to be Majorana particles.
+ Di (his basis. a complex svimmetric neutrinomass matrix can be diagonalized by a," In this basis, a complex symmetric neutrinomass matrix can be diagonalized by a"
+"A. A better choice for approximating the Band function is to require that the asymptotic behavior (for very large and very small 7.) of the broken power law matches the Band function, while leaving the break energy for the broken power law a tree parameter: The value of the break energy c, 1s given by the energy at which the two branches of the broken power law are equal to each other: Which results in an effective break energy that is independent of the spectral indices: where e is Euler's number.","A. A better choice for approximating the Band function is to require that the asymptotic behavior (for very large and very small $E_\gamma$ ) of the broken power law matches the Band function, while leaving the break energy for the broken power law a free parameter: The value of the break energy $\bar{\epsilon}_\gamma$ is given by the energy at which the two branches of the broken power law are equal to each other: Which results in an effective break energy that is independent of the spectral indices: where $e$ is Euler's number."
+ Figure | also shows approximation B with the same parameters as before., Figure \ref{fig:bandA} also shows approximation B with the same parameters as before.
+" Given a neutrino spectrum d«V,E, the expected number of events in a neutrino telescope where .1 is the muon effective area (1. Km? for IceCube/KM3NET). D,(E,:E"") is the probability of a neutrino of energy 77, to produce a muon withenergy equal or greater than the neutrino telescope threshold ο (we assume 100 GeV) and S(£,,.@) is Earth's attenuation."," Given a neutrino spectrum $dN_\nu/dE_\nu$, the expected number of events in a neutrino telescope is: where $A^{\mu}$ is the muon effective area (1 $^2$ for IceCube/KM3NET), $P_\mu(E_\nu;E^{\mathrm{min}}_\mu)$ is the probability of a neutrino of energy $E_\nu$ to produce a muon withenergy equal or greater than the neutrino telescope threshold $E^{\mathrm{min}}_\mu$ (we assume 100 GeV) and $S(E_\nu,\theta)$ is Earth's attenuation."
+"The probability ευ:ο) is given by: where <R,CE,:Eyin)> is the average muon range given a neutrino energy Z7, and a muon threshold i 77"".",The probability $P_\mu(E_\nu;E^{\mathrm{min}}_\mu)$ is given by: where $<R_\mu(E_\nu;E_\mu^{\mathrm{min}})>$ is the average muon range given a neutrino energy $E_\nu$ and a muon threshold $E_\mu^{\mathrm{min}}$ .
+ Earth's attenuation factor is f=given by:, Earth's attenuation factor is given by:
+temperature profiles are from Sandersetal.(2010).,temperature profiles are from \citet{sfsp10}.
+. The data points for the predicted quantities use |£|—0.75. but we also show upper and lower limits calculated by setting [£|=0.5 (dashed lines) and £|=1 (dotted lines). respectively.," The data points for the predicted quantities use $|\xi | = 0.75$, but we also show upper and lower limits calculated by setting $|\xi | = 0.5$ (dashed lines) and $|\xi | = 1$ (dotted lines), respectively."
+ Equations. (29)) (36)) do not imply a specifie causal relationship between the five quantities no. 7. D. (yay. and £L.," Equations \ref{eqn:Bprofile}) ) – \ref{eqn:kprofile}) ) do not imply a specific causal relationship between the five quantities $n_{\rm e}$, $T$, $B$, $U_{\rm rms}$ and $L$."
+ Given any two quantities. our theory can predict the other three.," Given any two quantities, our theory can predict the other three."
+ For example. if observations determined the profiles of n. and Cj. for a given cluster. rather than of Εν and 7. then our theory would predict 7. 5 and L.," For example, if observations determined the profiles of $n_{\rm e}$ and $U_{\rm rms}$ for a given cluster, rather than of $n_{\rm e}$ and $T$, then our theory would predict $T$, $B$ and $L$."
+ Ideally. our theory would be put to the test if three or more of these profiles were known observationally.," Ideally, our theory would be put to the test if three or more of these profiles were known observationally."
+ The predicted value for > near the centre of the core is =15 — 2]pi. decreasing to =5 — 7µία at the outer core boundary.," The predicted value for $B$ near the centre of the core is $\simeq 15$ – $21~\mu{\rm G}$, decreasing to $\simeq 5$ – $7~\mu{\rm G}$ at the outer core boundary."
+ To our knowledge. observational estimates of the magnetic-field strength in A1835 have not yet appeared in the literature.," To our knowledge, observational estimates of the magnetic-field strength in A1835 have not yet appeared in the literature."
+ As a radio mini-halo has recently been detected in AI835 (Murgiaetal. 2009).. there is hope for a magnetic field measurement there. although this might be quite a difficult task for such a distant cluster.," As a radio mini-halo has recently been detected in A1835 \citep{mgmfgtc09}, there is hope for a magnetic field measurement there, although this might be quite a difficult task for such a distant cluster."
+" The predicted turbulent velocity dispersion C4.10 — 270kms+ throughout the core. attaining a value C,114 — 162kms Lata radius of 30kpe."," The predicted turbulent velocity dispersion $U_{\rm rms}\sim 70$ – $270~{\rm km~s}^{-1}$ throughout the core, attaining a value $U_{\rm rms}\simeq 114$ – $162~{\rm km~s}^{-1}$ at a radius of $30~{\rm kpc}$."
+ This is within the 182kmsga+ upper limit obtained by Sandersetal.(2010). by measuring emission lines within 230 kpc. which is denoted on the plot by a thick arrow.," This is within the $182~{\rm km~s}^{-1}$ upper limit obtained by \citet{sfsp10} by measuring emission lines within $r\simeq 30~{\rm kpc}$ , which is denoted on the plot by a thick arrow."
+ A more conservative observational estimate of UnsS274kms+ was derived by treating the cluster as a point source and not applying any spatial smoothing (Sandersetal. 2010)., A more conservative observational estimate of $U_{\rm rms}\lesssim 274~{\rm km~s}^{-1}$ was derived by treating the cluster as a point source and not applying any spatial smoothing \citep{sfsp10}.
+. The predicted turbulence scale is £L.~0.2 — 0.7kpe near the centre of the core. increasing outwards to ~70 — 200kpc near the temperature maximum.," The predicted turbulence scale is $L\sim 0.2$ – $0.7~{\rm kpc}$ near the centre of the core, increasing outwards to $\sim 70$ – $200~{\rm kpc}$ near the temperature maximum."
+ The predicted diffusion coefticient Arun Fises sharply from —1077em?s* near the core centre to ~JOen?s “atthe outer core boundary., The predicted diffusion coefficient $\kappa_{\rm turb}$ rises sharply from $\sim 10^{28}~{\rm cm^2~s}^{-1}$ near the core centre to $\sim 10^{31}~{\rm cm^2~s}^{-1}$ at the outer core boundary.
+" Given the uncertainties discussed in Sections ?? and ??.. itis encouraging that these values are comparable to the inferred diffusion coefficients ~107"" 10%em?&7 in a variety of observed clusters (Rebuscoetal.2005.2006.2008:David&Nulsen 2008)."," Given the uncertainties discussed in Sections \ref{sec:velocities} and \ref{sec:scales}, it is encouraging that these values are comparable to the inferred diffusion coefficients $\sim 10^{29}$ -- $10^{30}~{\rm cm}^2~{\rm s}^{-1}$ in a variety of observed clusters \citep{rcbf05,rcbf06,rcsbf08,dn08}."
+. The plasma microphysics responsible for parallel viscous heating is in no way unique to cool-core clusters., The plasma microphysics responsible for parallel viscous heating is in no way unique to cool-core clusters.
+ Since plenty of turbulence is expected to be present (stirred. by mergers. etc.).," Since plenty of turbulence is expected to be present (stirred by mergers, etc.),"
+ pressure anisotropies will develop and will presumably be maintained at a marginal level as described in Section ??.., pressure anisotropies will develop and will presumably be maintained at a marginal level as described in Section \ref{sec:heat}.
+ While heating and cooling times are quite long in such clusters and the situation can be quite time-dependent. we would nevertheless like to explore what the conjecture of an approximate local heating- balance would imply. with the caveat that such a balance may in principle take a very long time to be established.," While heating and cooling times are quite long in such clusters and the situation can be quite time-dependent, we would nevertheless like to explore what the conjecture of an approximate local heating-cooling balance would imply, with the caveat that such a balance may in principle take a very long time to be established."
+ We will see that such a balance leads to quite reasonable predictions that seem to be borne out by observational data., We will see that such a balance leads to quite reasonable predictions that seem to be borne out by observational data.
+ In isothermal clusters. equations (29)) — (353) imply that vThe sealing: Dxης172 innon-cool-core clusters has in fact already been observationally inferred for A2382 (Guidettietal.2008 ).. Coma (Bonafedeetal.2010) and A665 (Vaccaetal.2010).," In isothermal clusters, equations \ref{eqn:Bprofile}) ) – \ref{eqn:Lprofile}) ) imply that The scaling $B\propto n_{\rm e}^{1/2}$ innon-cool-core clusters has in fact already been observationally inferred for A2382 \citep{gmgpgrcf08}, , Coma \citep{bfmggddt10} and A665 \citep{vmgfgob10}."
+". The central density n...25.10Ecm"" and temperature Tz2.9keV in A2382 (Ebelingetal.1996). implies a thermal-equilibrium magnetic-field strength. Boo3.1£07pC. in excellent agreement with the ~3yi estimate derived from rotation measure observations (Guidettietal.2008)."," The central density $n_{\rm e,c}\simeq 5\times 10^{-3}~{\rm cm}^{-3}$ and temperature $T\simeq 2.9~{\rm keV}$ in A2382 \citep{evb96} implies a thermal-equilibrium magnetic-field strength $B_{\rm c}\simeq 3.1~\xi^{-1/2}~\mu{\rm G}$, in excellent agreement with the $\sim 3~\mu{\rm G}$ estimate derived from rotation measure observations \citep{gmgpgrcf08}."
+". For A2255. Dos022.107em (CFerettietαἱ.1997:Govonial.2006) and 7=3.5keV (Davis&White1998.however.seeSakelliou who find temperature variations across A2255 from 7—5.5 — 8.5keV). the thermal-equilibrium magnetic-tield strength D,22:2£.L7pC."," For A2255, $n_{\rm e,c}\simeq 2\times 10^{-3}~{\rm cm}^{-3}$ \citep{fbgn97,gmfgdt06} and $T\simeq 3.5~{\rm keV}$ \citep[however, see \citealt{sp06}, , who find temperature variations across A2255 from $T\sim 5.5$ – $8.5~{\rm keV}$, the thermal-equilibrium magnetic-field strength $B_{\rm c}\simeq 2.2~\xi^{-1/2}~\mu{\rm G}$."
+ This compares favourably with the observational estimate 5.—2.5µία obtained by Govoni who used Dxni7 in their analysis.," This compares favourably with the observational estimate $B\sim 2.5~\mu{\rm G}$ obtained by \citet{gmfgdt06}, who used $B\propto n_{\rm e}^{1/2}$ in their analysis."
+ In Table |. we list these and other central magnetic-tield strength predictions.," In Table \ref{tab:bfields}, we list these and other central magnetic-field strength predictions."
+ Very recently there have been observational estimates of the magnetic-field strength profile in the non-cool-core Coma cluster by Bonafedeetal.(2010)., Very recently there have been observational estimates of the magnetic-field strength profile in the non-cool-core Coma cluster by \citet{bfmggddt10}.
+. They found that the best-fitting profile of the magnetic-field strength is with a 7;-model (Cavaliere&Fusco-Femiano1976) electron density profile where ny=3.4410%em Fore =291kpe and 3= 0.75.," They found that the best-fitting profile of the magnetic-field strength is with a $\beta$ -model' \citep{cf76} electron density profile where $n_0 = 3.44 \times 10^{-3}~{\rm cm}^{-3}$ , $r_{\rm c} = 291~{\rm kpc}$ and $\beta = 0.75$ ."
+ Taking Coma to be an isothermal cluster with temperature 8.2keV. (Arnaudetal.2001)... we find that the implied parallel viscous heating rate is (from eq.," Taking Coma to be an isothermal cluster with temperature $8.2~{\rm keV}$ \citep{arnaud01}, we find that the implied parallel viscous heating rate is (from eq."
+ I+ using eq. 4310) , \ref{eqn:heating3} using eq. \ref{eqn:bcoma}) )
+By way of comparison. the implied Bremsstrahlung cooling rate is (from eq. 30) ," By way of comparison, the implied Bremsstrahlung cooling rate is (from eq. \ref{eqn:cooling2}) )"
+"where n, is given by equation (42)).", where $n_{\rm e}$ is given by equation \ref{eqn:nefit}) ).
+ While it is rather curious that these two rates are not only comparable. but also have the same radial sealing. we are not on solid ground applying our estimates to Coma because the cooling time for it implied by equation (44)) is losaonfPjQ~ the Hubble time at the cluster centre. increasing outwards in radius.," While it is rather curious that these two rates are not only comparable, but also have the same radial scaling, we are not on solid ground applying our estimates to Coma because the cooling time for it implied by equation \ref{eqn:comacooling}) ) is $t_{\rm cool} \sim n T / Q^- \sim$ the Hubble time at the cluster centre, increasing outwards in radius."
+ The other isothermal clusters mentioned above and in Table | have relatively shorter cooling times. so one could in principle have a heating-cooling balance — but no radial dependence of the magnetic-field strength has so far been measured.," The other isothermal clusters mentioned above and in Table \ref{tab:bfields} have relatively shorter cooling times, so one could in principle have a heating-cooling balance – but no radial dependence of the magnetic-field strength has so far been measured."
+ Bearing in mind that thermal conduction models routinely fail in the innermost regions of cool cluster cores (e.g.Zakamska& 2004). it isimportant to note that parallel viscous δηης should be especially important in these relatively cold (7~ keV) and strongly-magnetised(D~10 7G) regions.," Bearing in mind that thermal conduction models routinely fail in the innermost regions of cool cluster cores \citep[e.g.][]{zn03,markevitch03,kaastra04,gmpd04}, , it isimportant to note that parallel viscous heating should be especially important in these relatively cold $T\sim 1~{\rm keV}$ ) and strongly-magnetised$B\sim 10~\mu{\rm G}$ ) regions."
+ A balance between parallel viscous heatingand radiative cooling. however. does contain the implicit assumption that the thermal conduction is relativelyunimportant.," A balance between parallel viscous heatingand radiative cooling, however, does contain the implicit assumption that the thermal conduction is relativelyunimportant."
+ This ought to be checked., This ought to be checked.
+ The Spitzer(1962) electron thermal diffusion coefficient is, The \citet{spitzer62} electron thermal diffusion coefficient is
+emission lines and has higher S/N. Also fibre cross-talk.,emission lines and has higher S/N. Also fibre cross-talk.
+ Radio source identifications of objects in the GadbGs-RASS catalogue were found. by crossmatching it with the l4GGllz NRAO VLA Sky Survey (NVSS: ?2)) at A and the MMlIz Sydney University Molonglo Sky Survey (SUMSS: ?2)) at 0<307., Radio source identifications of objects in the 6dFGS-RASS catalogue were found by crossmatching it with the GHz NRAO VLA Sky Survey (NVSS; \citet{nvss}) ) at $\delta > -40^{\circ}$ and the MHz Sydney University Molonglo Sky Survey (SUMSS; \citet{sumss}) ) at $\delta < -30^\circ$.
+ Candidate SUAISS and NVSS matches to 6dE€S objects were confirmed using the method outlined by 2. for the αςNVSS sample., Candidate SUMSS and NVSS matches to 6dFGS objects were confirmed using the method outlined by \citet{mauch} for the 6dFGS–NVSS sample.
+ The SUAISS anc NVSS radio source catalogues were searched for components within aarcesec of cach 6dECS position., The SUMSS and NVSS radio source catalogues were searched for components within arcsec of each 6dFGS position.
+ Those with a single radio component that were separated bv more than aarcsec [rom, Those with a single radio component that were separated by more than arcsec from
+and have subregions of high and low frec-free optical depth.,and have subregions of high and low free-free optical depth.
+ Their flux. during the accretion stage (while they Licker) is dominated by the denser. optically thicker (mp2 1) subregions. so their behaviour is closer to cq. (," Their flux during the accretion stage (while they flicker) is dominated by the denser, optically thicker $\tau_\mathrm{ff} > 1$ ) subregions, so their behaviour is closer to eq. ("
+4) than to eq. (,4) than to eq. (
+3).,3).
+ The on-line version of this paper contains a movie of zr; for Run Boas viewed from the Z-axis (line of sight perpendicular to the plane of the accretion Low)., The on-line version of this paper contains a movie of $\tau_{ff}$ for Run B as viewed from the Z-axis (line of sight perpendicular to the plane of the accretion flow).
+ The elumpiness ancl intermediate-to-LIarge optical depth of hese regions are also the reasons behind. their rising spectral indices up to relatively large frequencies (νο100 (11) without a significant contribution from dust. emission (sce the analytical discussions of Ignace Churchwell 2004 and Ixeto et al., The clumpiness and intermediate-to-large optical depth of these regions are also the reasons behind their rising spectral indices up to relatively large frequencies $\nu>100$ GHz) without a significant contribution from dust emission (see the analytical discussions of Ignace Churchwell 2004 and Keto et al.
+" 2008. for an analysis of these simulations see ""aper LL)."," 2008, for an analysis of these simulations see Paper II)."
+ As for the variability. the large optical depths cause the size ancl flux of the simulated regions to να well correlated with each other. anc anticorrelatec with he density of the central ionized. gas. (sec Section 3.6).," As for the variability, the large optical depths cause the size and flux of the simulated regions to be well correlated with each other, and anticorrelated with the density of the central ionized gas (see Section 3.6)."
+ The neutral aceretion flow in whieh the ionizing sources are embedded: is filamentary ancl prone. to gravitational instability (further discussion is in Section 3 of Paper 1. see also Paper HD).," The neutral accretion flow in which the ionizing sources are embedded is filamentary and prone to gravitational instability (further discussion is in Section 3 of Paper III, see also Paper II)."
+ The changes in the density of the regions are à consequence of their passage through density enhancements in the quickly evolving accretion Low., The changes in the density of the regions are a consequence of their passage through density enhancements in the quickly evolving accretion flow.
+ Figure 1 shows the elobal temporal evolution of the region in Run A (single sink particle)., Figure 1 shows the global temporal evolution of the region in Run A (single sink particle).
+" The 2-en flux 555) observed from orthogonal directions and the mass of 1e ionizing star (M,) are plotted against time.", The 2-cm flux $S_{\rm 2cm}$ ) observed from orthogonal directions and the mass of the ionizing star $M_\star$ ) are plotted against time.
+ The global emporal trend of the region is to expand and become xighter., The global temporal trend of the region is to expand and become brighter.
+ However. fast temporal variations are seen at all Ίο stages of the evolution.," However, fast temporal variations are seen at all the stages of the evolution."
+ The fluxes in the projections dong the three dillerent. cartesian axes follow cach other ‘losely., The fluxes in the projections along the three different cartesian axes follow each other closely.
+ For the rest of the analysis. the Z-axis projection. a ine of sight perpendicular to the plane of the accretion Low. is usec.," For the rest of the analysis, the Z-axis projection, a line of sight perpendicular to the plane of the accretion flow, is used."
+" The region is always faint (Sao,<1 Jv at the assumed distance of 2.65 kpe) for M,<25M..", The region is always faint $S_{\rm 2cm}<1$ Jy at the assumed distance of 2.65 kpc) for $M_\star<25~\Msun$.
+ Past this point. the region is brighter than 1 Jy 86% of the time (Fig.," Past this point, the region is brighter than 1 Jy 86 of the time (Fig."
+ 1)., 1).
+ A similar analysis of the region around the most massive star in Run B (multiple sinks) is shown in Fig., A similar analysis of the region around the most massive star in Run B (multiple sinks) is shown in Fig.
+ 2., 2.
+ Only the Z-axis projection. te. a line of sight perpendicular to plane of the aceretion Low is used. since only from. this viewing angle the brightest region is well separated [rom other regions at all times (the Dux movies are presented in Paper D).," Only the Z-axis projection, i.e., a line of sight perpendicular to plane of the accretion flow is used, since only from this viewing angle the brightest region is well separated from other regions at all times (the flux movies are presented in Paper I)."
+" The extra fragmentation in Run B translates into a weaker accretion How and a lower-miass ionizing star as compared to that of Run A. a process referred. to as ""Tragmoentation-induced starvation! (à theoretical discussion of this process in presented in Paper HE)."," The extra fragmentation in Run B translates into a weaker accretion flow and a lower-mass ionizing star as compared to that of Run A, a process referred to as 'fragmentation-induced starvation' (a theoretical discussion of this process in presented in Paper III)."
+ Pherefore. the xightest region in Run D (Fig.," Therefore, the brightest region in Run B (Fig."
+ 2) is weaker than the region in Run A (Eig., 2) is weaker than the region in Run A (Fig.
+ 1)., 1).
+ Figures 1 and 2 do not show times ater than /=100 kvr in order to facilitate their comparison., Figures 1 and 2 do not show times later than $t=100$ kyr in order to facilitate their comparison.
+ Both runs continue past this time. but in Runa D accretion onto the most massive star stops at /=LOO kyr. while in tun A the artificial suppression of the fragmentation leacks o an unrealistically large mass for the ionizing star at later imes (see Paper 1).," Both runs continue past this time, but in Run B accretion onto the most massive star stops at $t=109$ kyr, while in Run A the artificial suppression of the fragmentation leads to an unrealistically large mass for the ionizing star at later times (see Paper I)."
+ regions are highly variable both in Run A and tun D. However. since the suppression of fragmentation in Run X produces a larger accretion [low and a most massive tonizine star. Run A presents larger [lux variations han Run D. Figure 3 shows the flux changes over the evolution of the regions in both runs.," regions are highly variable both in Run A and Run B. However, since the suppression of fragmentation in Run A produces a larger accretion flow and a most massive ionizing star, Run A presents larger flux variations than Run B. Figure 3 shows the flux changes over the evolution of the regions in both runs."
+ A comparison of the fractional variations of the regions shows that.," A comparison of the fractional variations of the regions shows that,"
+As an incidental remark. the existence of the long range ticle - void alignment tells us that the void anisotropy is not a mere artefact of the NVE method chopping spherical voids in two (which would of course lead to non-spherical voids having correlated alignments),"As an incidental remark, the existence of the long range tide - void alignment tells us that the void anisotropy is not a mere artefact of the WVF method chopping spherical voids in two (which would of course lead to non-spherical voids having correlated alignments)."
+ We have investigated the shapes ancl alignments of voids found in the large scale. matter distribution of the VLS CIP N-body simulation of structure formation in a .MCDM universe., We have investigated the shapes and alignments of voids found in the large scale matter distribution of the VLS GIF N-body simulation of structure formation in a $\Lambda$ CDM universe.
+ The void sample has been identified: using the Watershed. Void. Finder (WVE) technique., The void sample has been identified using the Watershed Void Finder (WVF) technique.
+ Although voids tend. to. be less fattened or elongated than the halos in the dark matter distribution. they are nevertheless quite nonspherical: they are slightly: prolate with axis ratios on the order of e:bαc0.50.7: 1.," Although voids tend to be less flattened or elongated than the halos in the dark matter distribution, they are nevertheless quite nonspherical: they are slightly prolate with axis ratios on the order of $c:b:a\approx 0.5:0.7:1$ ."
+ Two important factors contribute to this Hattening., Two important factors contribute to this flattening.
+ Internally. voids tend to become more spherical as they expand (Ieke1984).," Internally, voids tend to become more spherical as they expand \citep{icke1984}."
+. However. they will never be able to reach perfect sphericity before meeting up with surrounding structures.," However, they will never be able to reach perfect sphericity before meeting up with surrounding structures."
+ Furthermore. external tidal forces will contribute substantially to the anisotropic developmoent of the voids.," Furthermore, external tidal forces will contribute substantially to the anisotropic development of the voids."
+ In this πεν we have established the inlluence of the tidal fields as being an important agent in the dynamical evolution of voids., In this study we have established the influence of the tidal fields as being an important agent in the dynamical evolution of voids.
+ Indeed. our study. provides strong confirmation that tidal forces dominate the shaping of the Cosmic Web. as emphasized long ago by Bondοἱa£(1996).," Indeed, our study provides strong confirmation that tidal forces dominate the shaping of the Cosmic Web, as emphasized long ago by \cite{bondweb1996}."
+. We have also investigated the relative orientation of the voids., We have also investigated the relative orientation of the voids.
+ “Phe orientation of voids appears to be strongly correlated with alignments spanning clistances 230h.IMpe., The orientation of voids appears to be strongly correlated with alignments spanning distances $>30h^{-1}\hbox{Mpc}$.
+ This is true of both the shortest and longest. axes of the void shape., This is true of both the shortest and longest axes of the void shape.
+ This coherence is quite conspicuous in plots of he density distribution., This coherence is quite conspicuous in plots of the density distribution.
+ To test for such void alignments in observational data a follow-up study will apply the WVE ormalism to galaxy redshift surveys., To test for such void alignments in observational data a follow-up study will apply the WVF formalism to galaxy redshift surveys.
+ We find an intimate link between the cosmic tidal ielel and the void orientations., We find an intimate link between the cosmic tidal field and the void orientations.
+ Over a range of scales we find a strong alignment of the voids with the tidal ield arising [rom the smoothed density distribution., Over a range of scales we find a strong alignment of the voids with the tidal field arising from the smoothed density distribution.
+ Given hat the alignment correlations remain significant on scales considerably exeeeding twice the tvpical void size. our results show that the long range external field is responsible or the alignment of the voids.," Given that the alignment correlations remain significant on scales considerably exceeding twice the typical void size, our results show that the long range external field is responsible for the alignment of the voids."
+ This confirms the view that he large scale tidal force field is the main agent for the spatial organization of the Cosmic Web., This confirms the view that the large scale tidal force field is the main agent for the spatial organization of the Cosmic Web.
+ We have argued that the large scale tidal field is not only the main agent responsible for the shapes of the voids but also for the resulting alignments of the voids., We have argued that the large scale tidal field is not only the main agent responsible for the shapes of the voids but also for the resulting alignments of the voids.
+ Locally. the orientation of a void turns out to be strongly aligned. with the tidal force field generated by structures on scales up to atleast 20.305.1Mpc.," Locally, the orientation of a void turns out to be strongly aligned with the tidal force field generated by structures on scales up to at least $20-30h^{-1}\hbox{\rm Mpc}$."
+ On scales comparable to the average void size we also found indications of nonlinear effects. with the void's orientation reacting to the small-scale nonlinear inlluences.," On scales comparable to the average void size we also found indications of nonlinear effects, with the void's orientation reacting to the small-scale nonlinear influences."
+ This conclusion. agrees with that of à similar. study of halo alignments by Leeetal.(2007)., This conclusion agrees with that of a similar study of halo alignments by \cite{leespringel2007}.
+.. However. the alignment of halos tends to be strongly attenuated by local nonlinear effects (e.g.vanHaarlem&deWeveacrt1993:Aragón-Calvoetal.2007) rendering the final halo alignment signal weaker than that for voids.," However, the alignment of halos tends to be strongly attenuated by local nonlinear effects \citep[e.g.][]{haarwey1993, aragon2007} rendering the final halo alignment signal weaker than that for voids."
+ One aspect we have not vet adressed. in detail is the entanglement of initial conditions and the influence of the tical Geld., One aspect we have not yet adressed in detail is the entanglement of initial conditions and the influence of the tidal field.
+ On Large scales. where the density Lluctuations are in the lincar regime. these are different aspects of the inhomogeneous matter distribution: if one is known the other can be determined.," On large scales, where the density fluctuations are in the linear regime, these are different aspects of the inhomogeneous matter distribution: if one is known the other can be determined."
+ In this respect the indication of nonlinear effects in the tical-void alignment νο (lig. 4)), In this respect the indication of nonlinear effects in the tidal-void alignment ${\mathcal A}_{TS}$ (fig. \ref{fig:cor2}) )
+ at small scales is interesting., at small scales is interesting.
+ They may very well be an expression of the strong influence of coherent overdense filaments and. dense compact clusters on the evolution of small voids in the outer regions of large voids., They may very well be an expression of the strong influence of coherent overdense filaments and dense compact clusters on the evolution of small voids in the outer regions of large voids.
+ vandeWeveaertefaf.(2004) argued that the collapse of small voids. an essential aspect of the hierarchical evolution of voids (Sheth&vandeWeveacrt2004).. manifests itself in a tidally induced anisotropic contraction of small uncdercensitics at the void. bouncaries.," \cite{weysheth2004} argued that the collapse of small voids, an essential aspect of the hierarchical evolution of voids \citep{shethwey2004}, manifests itself in a tidally induced anisotropic contraction of small underdensities at the void boundaries."
+ The dvnanmies and the fate of these small unclerclensities is often clecisively inlluenced. by external anisotropic forces (vancleWeveacrt&Babul 1996)., The dynamics and the fate of these small underdensities is often decisively influenced by external anisotropic forces \citep{weybabul1996}.
+. While such forces may allect the nonlinear evolution of halos and clusters to a considerable extent (Bond&Myers1996:Shethetal.2001).. they do not eo as [ar as deciding their fate.," While such forces may affect the nonlinear evolution of halos and clusters to a considerable extent \citep{bondmyers1996,sheth2001}, they do not go as far as deciding their fate."
+ This work is part of a study of the evolution of voids under the influence of external inlluences., This work is part of a study of the evolution of voids under the influence of external influences.
+ |. proper framework for following the gradual hierarchical buildup of void regions under the influence of external tidal forces is provided. by the void-patch. description. (Platen.vandeWevgaert&Jones2008:Bond.efa£ 1996).," A proper framework for following the gradual hierarchical buildup of void regions under the influence of external tidal forces is provided by the void-patch description \citep{platen2008, bondweb1996}."
+. Fo some extent 16 void-patch formalism is à more accurate approximation X reality than for the equivalent. peak-patch: because. of jer approximate uniformity void evolution is accurately escribed by the ellipsoidal moclel., To some extent the void-patch formalism is a more accurate approximation of reality than for the equivalent peak-patch: because of their approximate uniformity void evolution is accurately described by the ellipsoidal model.
+ A perhaps even more complex issue is that of. the --dentitv of the outer regions of voids. the regions where 1ο void merges into the surrounding matter distribution.," A perhaps even more complex issue is that of the identity of the outer regions of voids, the regions where the void merges into the surrounding matter distribution."
+ _Ve have argued that this will play an important role in etermining the shape of voids., We have argued that this will play an important role in determining the shape of voids.
+ However. a proper definition ofa void boundary does not (vet) exist.," However, a proper definition of a void boundary does not (yet) exist."
+ Areuably the WVE has proven to represent a major advance along these lines ane we should soon be able to be more definite on the issue., Arguably the WVF has proven to represent a major advance along these lines and we should soon be able to be more definite on the issue.
+ llaving established the external influcnees on a void's evolution. we may try to understand. the sensitivity of the void. population to the underlying cosmology.," Having established the external influences on a void's evolution, we may try to understand the sensitivity of the void population to the underlying cosmology."
+ This depends to a considerable degree on the dillerences in scale dependence of the tidal force fields., This depends to a considerable degree on the differences in scale dependence of the tidal force fields.
+ Lee&Park(2007). did find a considerable inlluence., \cite{leepark2007} did find a considerable influence.
+ Vhis will open the door to the use of voids in determining the global cosmology., This will open the door to the use of voids in determining the global cosmology.
+ We eratefully ackowledge the use of the Virgo Consortium simulation., We gratefully ackowledge the use of the Virgo Consortium simulation.
+ We thank Miguel Aragónn-C'alvo. [ου useful discussions on various technical aspects of this paper., We thank Miguel Aragónn-Calvo for useful discussions on various technical aspects of this paper.
+arguments. the inference about simall-scale radio ciission still stands.,"arguments, the inference about small-scale radio emission still stands."
+ The detection of two point sources at radio frequeucies also raises the possibility that the eudssiou origmates from two quasars apart. or at a projected separation of 5.1 kpc.," The detection of two point sources at radio frequencies also raises the possibility that the emission originates from two quasars apart, or at a projected separation of 5.1 kpc."
+ Borosou&Lauer(2009) noted that the probability for such a raudom projection in their sample is 0.0032., \citet{bor09} noted that the probability for such a random projection in their sample is 0.0032.
+ Therefore. the two quasars are inost likely not due to a random projection. but are likely physically related. ic. a binary quasar svstem.," Therefore, the two quasars are most likely not due to a random projection, but are likely physically related, i.e. a binary quasar system."
+ A remarkable radio-loud case of a binary svsteni of active galactic nuclei CACUNS) is seen in 3€775. where two systems of two-sided jets cluanate from two close point sources with a projected separation of ~7.5 kpe (Owenetal.1985).," A remarkable radio-loud case of a binary system of active galactic nuclei (AGNs) is seen in 75, where two systems of two-sided jets emanate from two close point sources with a projected separation of $\sim 7.5$ kpc \citep{owe85}."
+. Compact binary ACNs were also revealed with observations of nearby. svsteimis., Compact binary AGNs were also revealed with observations of nearby systems.
+ Particularly clear cases are 66210 with a ~1 kpe separation (Ixoniossaetal.2003).. and. £163. with ao3.8 kpc separation (Bianchictal.2008).," Particularly clear cases are 6240 with a $\sim 1$ kpc separation \citep{kom03}, and 463 with a $\sim 3.8$ kpc separation \citep{bia08}."
+.. Thus. SDSS J1536|OLLI may be another example of a compact binary ACN system.," Thus, SDSS J1536+0441 may be another example of a compact binary AGN system."
+ A systematic study of the abunudauce of binary quasars in the SDSS was carried out by Ieunawietal.(2006). who found a projected correlation function of the form (RoopfO.LAIpeh1jL5 on sales of 10 kpce in proper length (proper length is used given the abseuce of a Dibble Sow ou these small scales).," A systematic study of the abundance of binary quasars in the SDSS was carried out by \citet{hen06}, who found a projected correlation function of the form $(R_{\rm prop}/0.43{\rm Mpc}{\rm
+h}^{-1})^{-1.48}$ on scales of 10–40 kpc in proper length (proper length is used given the absence of a Hubble flow on these small scales)."
+ Extrapolating to 5 kpc (using 5h= 0.71) we ect a correlation fiction of 1175. ie. an observed surface deusity 1175 larger han expected for random projection.," Extrapolating to 5 kpc (using $h=0.71$ ) we get a correlation function of 1175, i.e. an observed surface density 1175 larger than expected for random projection."
+" Using the raudom-xojection estimate of 0.0032 (Boroson&Lauer2009).. he expected observed umber is actually 3.76. Ίο, 3 to L such binaries are expected."," Using the random-projection estimate of 0.0032 \citep{bor09}, the expected observed number is actually 3.76, i.e. 3 to 4 such binaries are expected."
+ The study of Uemnawi ct al., The study of Hennawi et al.
+ is owed ou quasars With aiean z~1.5 due to the smaller inber of quasars at lower 2. aud it does not extend down to 5 kpe due to the fiber angular resolution limit.," is based on quasars with a mean $z\sim 1.5$ due to the smaller number of quasars at lower $z$, and it does not extend down to 5 kpc due to the fiber angular resolution limit."
+ However. the study of Tennawi et al.," However, the study of Hennawi et al."
+ indicates that the xobabilitv to find a 5 kpe binary quasar in the sample uxed by Boroson&Lauer(2009) is ofthe order of nuit., indicates that the probability to find a 5 kpc binary quasar in the sample used by \citet{bor09} is of the order of unity.
+ Each of the two quasars in the binary nav be contributing its broad-line euissiou system to the total ieht. producing the double broad-line svsteii discovered x Boroson&Lauer(2009).," Each of the two quasars in the binary may be contributing its broad-line emission system to the total light, producing the double broad-line system discovered by \citet{bor09}."
+. Thus. rather than haviug a binary black-hole syste on a scale of 0.1-pc. we may lave a binary quasar svete on a 5-kpe scale. iiost likely residing within two stronely interacting ealaxies.," Thus, rather than having a binary black-hole system on a scale of 0.1-pc, we may have a binary quasar system on a 5-kpc scale, most likely residing within two strongly interacting galaxies."
+ The velocity separation of 3500 kia + (Doroson&Lauer2009) is rather huge. but not implausible iu a cluster of galaxies.," The velocity separation of 3500 km $^{-1}$ \citep{bor09} is rather large, but not implausible in a cluster of galaxies."
+ About half the clusters studied by Carlbereetal.(1996) show such an exteut of velocities., About half the clusters studied by \citet{car96} show such an extent of velocities.
+ The maxima velocity differeuces in the quasar binaries studied by. Heuuawietal.(2006). is ISTO kins 1. but that study imposed a cap of 2000 kins + on the binary velocity separation.," The maximum velocity differences in the quasar binaries studied by \citet{hen06} is 1870 km $^{-1}$, but that study imposed a cap of 2000 km $^{-1}$ on the binary velocity separation."
+ At a projected relative velocity of 3500 haus +. two galaxies cannot form a bound system. so the term binary quasar here docs not refer to a physically bouud binary svsteni.," At a projected relative velocity of 3500 km $^{-1}$, two galaxies cannot form a bound system, so the term binary quasar here does not refer to a physically bound binary system."
+ For some simall N-rav and radio selected samples of radio-quict AGNs. a link between the radio aud N-ray enission has been sugeested by the correlation between the radio aud the N-ray luminosities (Brinkiaunetal.2000:Salvatoetal.2001:Wane2006).," For some small X-ray and radio selected samples of radio-quiet AGNs, a link between the radio and X-ray emission has been suggested by the correlation between the radio and the X-ray luminosities \citep{bri00,sal04,wan06}."
+. While intriguing. such a radio/N-rav connection should be verified by using an unbiased survey.," While intriguing, such a radio/X-ray connection should be verified by using an unbiased survey."
+ The PG quasar siuuple is a complete optically-sclected sample (Borosou&Cacen1992). making it independent of radio aud N-rav biases.," The PG quasar sample is a complete optically-selected sample \citep{bor92}, making it independent of radio and X-ray biases."
+" Laor&Behar(2008) have used these radio-quiet PC quasars to demonstrate that (a) the radio and the N-rav luminosities are correlated over à larec range of ACN luminosity aud (0) the correlation follows EpiLy~1407. the well-established correlation for coronallv active cool stars (Caredel&Benz1993).. where Lpa=awl, ath GIIz and Ly is in the 0.2-20 keV. baud."," \citet{lao08} have used these radio-quiet PG quasars to demonstrate that (a) the radio and the X-ray luminosities are correlated over a large range of AGN luminosity and (b) the correlation follows $L_R / L_X \sim 10^{-5}$, the well-established correlation for coronally active cool stars \citep{gue93}, where $L_R = \nu L_\nu$ at 5 GHz and $L_X$ is in the 0.2-20 keV band."
+ For cool stars. the Ly/~107 relation is accepted as a manifestation of coronal heatiug by euergetic electrous following maeuetic reconnection. that subsequeutlv eives vise to X-ray cussion.," For cool stars, the $L_R / L_X \sim 10^{-5}$ relation is accepted as a manifestation of coronal heating by energetic electrons following magnetic reconnection, that subsequently gives rise to X-ray emission."
+ Dy analogy with cool stars. Laor&Behar(2008) conjecture that the radio emission in racdio- ACINS may also be related. to coronal maguetic activity.," By analogy with cool stars, \citet{lao08} conjecture that the radio emission in radio-quiet AGNs may also be related to coronal magnetic activity."
+ This coronal framework can be tested for SDSS J1536|OU by inakine use of its recent Swift observation o» Arzotmmanamianetal. (2009).. carried out im 2009 February [d and 5. just two weeks before our VLA observations.," This coronal framework can be tested for SDSS J1536+0441 by making use of its recent Swift observation by \citet{arz09}, , carried out in 2009 February 4 and 5, just two weeks before our VLA observations."
+ Arzommananuianetal.(2009) measure a 15-10 keV luminosity of 5«10! ere L1 with a spectral slope of 1.5.," \citet{arz09} measure a 0.5-10 keV luminosity of $5\times 10^{44}$ erg $^{-1}$, with a spectral slope of $-1.5$."
+" This gives pL,=2.3«10th ees Lat L keV. and thus Ly of ες10% ere 1 (seo Laor&Behar(2008) for the conversion of pL, at 1 keV to Ly)."," This gives $\nu L_\nu=2.3\times 10^{44}$ erg $^{-1}$ at 1 keV, and thus $L_X$ of $1.4\times 10^{45}$ erg $^{-1}$ (see \citet{lao08}
+ for the conversion of $\nu L_\nu$ at 1 keV to $L_X$ )."
+" The total radio huuinositv we find at 8.5 CGIIz is θες1079 ere 1, which extrapolates to Ly=8.3LO eres assundue a spectral slope of 0.0."," The total radio luminosity we find at 8.5 GHz is $6.4\times 10^{40}$ erg $^{-1}$, which extrapolates to $L_R=8.3\times 10^{40}$ erg $^{-1}$, assuming a spectral slope of $-0.5$."
+ We therefore ect Lg/Ly=5.9«10? which falls within the range of ratios seen by Laor&Behar(2008) for the radio-quiet PC quasars., We therefore get $L_R / L_X = 5.9\times 10^{-5}$ which falls within the range of ratios seen by \citet{lao08} for the radio-quiet PG quasars.
+" Thus. SDSS J1536|OLLL follows the radioταν relation of optically selected radio quiet quasars,"," Thus, SDSS J1536+0441 follows the radio/X-ray relation of optically selected radio quiet quasars."
+ The radio-quiet quasar SDSS οσο has two broad-line emission systems that Borosou&Lauer(2009) interpret as a candidate binary black-hole svsteu with a separation of 0.1 pe (0.02 mas)., The radio-quiet quasar SDSS J1536+0441 has two broad-line emission systems that \citet{bor09} interpret as a candidate binary black-hole system with a separation of 0.1 pc (0.02 mas).
+ Our new VLA imagine at 8.5 GIIz reveals two sources. separated by (5.1 kpc). within the quasars optical localization reeion.," Our new VLA imaging at 8.5 GHz reveals two sources, separated by (5.1 kpc), within the quasar's optical localization region."
+ Each radio source has a diameter of less than (1.9 kpc)., Each radio source has a diameter of less than (1.9 kpc).
+ Other radio-quiet quasars do exhibit double structures. sugeestiue that the radio double in SDSS 31536|0111 could be energized bw the candidate O.l-pc binary.," Other radio-quiet quasars do exhibit double structures, suggesting that the radio double in SDSS J1536+0441 could be energized by the candidate 0.1-pc binary."
+ Alternatively. the radio emission may arise from a binary system of quasars with a projected separation of 5.1 kpc. and those two quasars may be respousible for the two observed broad-line emission svstemis.," Alternatively, the radio emission may arise from a binary system of quasars with a projected separation of 5.1 kpc, and those two quasars may be responsible for the two observed broad-line emission systems."
+ Binary ACNs witli kpc-scale separations are known from radio and N-rav observations. and a few such system are expected in the Doroson&Lauer(2009). sample based ou the observed clustering of quasars down to a scaleof 10 kpe.," Binary AGNs with kpc-scale separations are known from radio and X-ray observations, and a few such system are expected in the \citet{bor09} sample based on the observed clustering of quasars down to a scaleof 10 kpc."
+refbars.,.
+. Given that excitation problems occur for some modes with frequencies above in the three well-known bright class members v Eri. LLac. and yPeg. it is quite encouraging to find models whose oscillation spectra are m full agreement with all the numerous observational constraints. except for the spectroscopic gravity and the excitation of the £23. pj mode.," Given that excitation problems occur for some modes with frequencies above in the three well-known bright class members $\nu\,$ Eri, Lac, and $\gamma\,$ Peg, it is quite encouraging to find models whose oscillation spectra are in full agreement with all the numerous observational constraints, except for the spectroscopic gravity and the excitation of the $\ell=3$, $_1$ mode."
+ In addition to the five remaining models. there are 38 models for which the mode excitation is fulfilled. except for the two highest frequencies in reffreqs.. r.e.. for one frequency less than the five models in the upper part of refmodels..," In addition to the five remaining models, there are 38 models for which the mode excitation is fulfilled, except for the two highest frequencies in \\ref{freqs}, i.e., for one frequency less than the five models in the upper part of \\ref{models}."
+" The evolutionary tracks to which these also ""acceptable"" models in the lower part of refmodels belong are indicated as thin dotted lines on refhrd.. for reference with respect to the five best models."," The evolutionary tracks to which these also “acceptable” models in the lower part of \\ref{models} belong are indicated as thin dotted lines on \\ref{hrd}, for reference with respect to the five best models."
+" These 43 models together cover a mass range of [10.5.12.0] MM... the full range X€[0.68.0.74] of the grid. Z€[0.014.0.018]. and cu, betweel zero and 0.2."," These 43 models together cover a mass range of $[10.5,12.0]$ $_\odot$, the full range $X\in [0.68, 0.74]$ of the grid, $Z\in
+[0.014, 0.018]$, and $\alpha_{\rm ov}$ between zero and 0.2."
+ That the initial. internal netallicity of the star seems to be somewhat higher (1c level) than the observed one at the surface for all the acceptable nodels. might have been introduced by the discrepancy in the gravity. which is not treated independently from. the nicrotubulence i à spectroscopic analysis.," That the initial internal metallicity of the star seems to be somewhat higher $\sigma$ level) than the observed one at the surface for all the acceptable models, might have been introduced by the discrepancy in the gravity, which is not treated independently from the microtubulence in a spectroscopic analysis."
+ It night also be due to effects of diffusion given 1180642’s longitudinal nagnetic field of up to GG (Hubrig et 22009. 2011).," It might also be due to effects of diffusion given 180642's longitudinal magnetic field of up to G (Hubrig et 2009, 2011)."
+ Almost all the seismic models have loge=3.83. which is ddex (2.507) above the best estimate from a classical spectroscopic analysis IID.," Almost all the seismic models have $\log\,g\simeq3.83$, which is dex $\sigma$ ) above the best estimate from a classical spectroscopic analysis II)."
+ We thus face another case where the seismic gravity differs from the spectroscopic gravity. just as for the B Cep pulsators € Oph (B21V. Briquet et 22007) and 446202 (O9V. Briquet et 22011).," We thus face another case where the seismic gravity differs from the spectroscopic gravity, just as for the $\beta\,$ Cep pulsators $\theta\,$ Oph (B2IV, Briquet et 2007) and 46202 (O9V, Briquet et 2011)."
+ Combining the seismic gravity with the equatorial surface rotation velocity of kkmss! and a radius of -6.8 Rs. leads to a surface rotation frequency of (rotation period of 13.3dd).," Combining the seismic gravity with the equatorial surface rotation velocity of $^{-1}$ and a radius of $R\simeq\,6.8\,$ $_\odot$, leads to a surface rotation frequency of (rotation period of d)."
+ We thus find that a factor of four occurs between the surface rotation frequency and the lowest frequency that was jointly detected in the CoRoT light curve and the spectroscopic time series., We thus find that a factor of four occurs between the surface rotation frequency and the lowest frequency that was jointly detected in the CoRoT light curve and the spectroscopic time series.
+ The frequency listed in reffreqgs is 12. times the rotation. frequency. within. the measurement errors.," The frequency listed in \\ref{freqs} is 12 times the rotation frequency, within the measurement errors."
+ The seismic rotation period that we obtained from our best-fitting models is 1n excellent agreement with the most probable period deduced from the magnetic field measurements., The seismic rotation period that we obtained from our best-fitting models is in excellent agreement with the most probable period deduced from the magnetic field measurements.
+ Continuing further with the models in refmodels.. we considered their entire oscillation spectrum in the range [0.2.40]d7!.. taking into account all rotational splittings according to the Ledoux formulation (e.g.. Aerts et 22010) for a rotation frequency of ed.. r.e.. assuming rigid rotation.," Continuing further with the models in \\ref{models}, we considered their entire oscillation spectrum in the range $[0.2,40]\,$, taking into account all rotational splittings according to the Ledoux formulation (e.g., Aerts et 2010) for a rotation frequency of , i.e., assuming rigid rotation."
+ In we show the mode inertia (defined as ((3.140) in Aerts et 22010) versus the oscillation frequencies for the seismic model with M=11.4 M.« listed in refmodels in the range [0.20].. for modes with £ 4.," In \\ref{inertia}, we show the mode inertia (defined as (3.140) in Aerts et 2010) versus the oscillation frequencies for the seismic model with $M=11.4\,$ $_\odot$ listed in \\ref{models} in the range $[0,20]\,$, for modes with $\ell=0,\dots,4$ ."
+ We note that the inertia follow a smooth curve as the radial order increases. except for some specific modes of £(=3 and 4 where the details of the trapping in the interior change owing to the interplay between the mode frequency and the Briinnt-Vinsailla and Lamb frequencies.," We note that the inertia follow a smooth curve as the radial order increases, except for some specific modes of $\ell=3$ and 4 where the details of the trapping in the interior change owing to the interplay between the mode frequency and the Brünnt-Väiisällä and Lamb frequencies."
+ The detected frequencies in the best fitting CoRoT light curve model deduced in PaperH are shown as dotted vertical lines. except for the harmonics of the dominant frequency.," The detected frequencies in the best fitting CoRoT light curve model deduced in I are shown as dotted vertical lines, except for the harmonics of the dominant frequency."
+ Full symbols denote modes predicted to be excited from the non-adiabatic computations done with the code MAD. while the open symbols are predicted to be stable.," Full symbols denote modes predicted to be excited from the non-adiabatic computations done with the code , while the open symbols are predicted to be stable."
+ All the combination frequencies (listed in. 22. of Paperll) in the range [4.8.5] can. be closely identified with frequencies of excited modes.," All the combination frequencies (listed in 2 of I) in the range $[4,8.5]\,$ can be closely identified with frequencies of excited modes."
+ This is also true for the four frequencies in the range [0.8.1.2].. which ean be explained by excited ¢ modes of intermediate to high radial order (7 typically between 10 and 30).," This is also true for the four frequencies in the range $[0.8,1.2]\,$, which can be explained by excited g modes of intermediate to high radial order $n$ typically between 10 and 30)."
+ All other combination frequencies found in Paperll can be attributed to model frequencies of non-excited p modes taking into account the observational and theoretical uncertainties., All other combination frequencies found in I can be attributed to model frequencies of non-excited p modes taking into account the observational and theoretical uncertainties.
+ If the frequency were produced by pulsation. it would have to correspond to an £(=3 mode IIL).," If the frequency were produced by pulsation, it would have to correspond to an $\ell=3$ mode II)."
+ There are indeed various (=3modes available. taking into account the rotational splitting (see the inset of refinertia)).," There are indeed various $\ell=3$modes available, taking into account the rotational splitting (see the inset of \\ref{inertia}) )."
+ These have radial orders from 53 to60., These have radial orders from 53 to60.
+ Given the, Given the
+considerations.,considerations.
+ Yet the departures from the energy equipartition between magnetic field ancl radiating electrons within the FRI lobes are never that large (INataoka&StawarzCrostonοἱal. 2005).. as demonstrated. also for several particular cases of GRGs (Ixonaral.2009.2011a.b).," Yet the departures from the energy equipartition between magnetic field and radiating electrons within the II lobes are never that large \citep{kat05,cro05}, as demonstrated also for several particular cases of GRGs \citep{kon09,iso09,iso11a,iso11b}."
+. Hence. in our analvsis we apply the standarcl mininunr-energv condition. but allow for a non-negligible contribution. of non-racliating parücles (ullvarelativistic or only mildlv-relativistic protons) to the total pressure inside the cocoons.," Hence, in our analysis we apply the standard minimum-energy condition, but allow for a non-negligible contribution of non-radiating particles (ultrarelativistic or only mildly-relativistic protons) to the total pressure inside the cocoons."
+ The ratio of the energv densitv stored in the lobes. magnetic field J? to that stored in all the lobes! particles is then Rey=(1+sinj)/[]4(A+1)]. where si;=2ainj+1 is the enerev slope of the electrons (wilh minimum and maxinnin Lorentz [actors αμ and μι. respectively) freshly accelerated at the jet termination shock. ancl A is the enerev/pressure fraction of the non-radiating particles.," The ratio of the energy density stored in the lobes' magnetic field $B$ to that stored in all the lobes' particles is then ${\cal R}_{\rm eq}=(1+s_{\rm inj})/[4\,(k^{\prime}+1)]$, where $s_{\rm inj}=2\,\alpha_{\rm inj}+1$ is the energy slope of the electrons (with minimum and maximum Lorentz factors $\gamma_{\rm min}$ and $\gamma_{\rm min}$, respectively) freshly accelerated at the jet termination shock, and $k^{\prime}$ is the energy/pressure fraction of the non-radiating particles."
+ The observational data summarized in 22 are fitted using the algorithm described in 33.1., The observational data summarized in 2 are fitted using the algorithm described in 3.1.
+ The results of the modeling are given in Table66., The results of the modeling are given in 6.
+" In the fitting procedure we [ix 1e following model parameters: adiabatic indices for (he plasma within and outside of the COCOODS «=j/ 5/3. minimum electron Lorentz factor 5,4,=1 and the masiniMeclron. Lorentz factor 544;=10*. as well as the inclination of the lobes 0.=90° (see assumption)."," In the fitting procedure we fix the following model parameters: adiabatic indices for the plasma within and outside of the cocoons $\hat{\gamma}_{\rm c} = \hat{\gamma}_{\rm IGM} = 5/3$ , minimum electron Lorentz factor $\gamma_{\rm min} = 1$ and the maximumelectron Lorentz factor $\gamma_{\rm max} = 10^7$, as well as the inclination of the lobes $\theta = 90\degr$ \citep[see][for more discussion regarding these model assumption]{mach07}."
+" In the anticipated density profile characterizing the ambient medium into which the lobes evolve. pir)=pytr/ag)"". we fix the characteristic radius of the density plateau e,= LOkkpe. and in (he case of the outer cocoons propagating through an undisturbed IGM the power-law index 3j=1.5."," In the anticipated density profile characterizing the ambient medium into which the lobes evolve, $\rho(r)=\rho_{0}\,(r/a_{0})^{-\beta}$, we fix the characteristic radius of the density plateau $a_0 = 10$ kpc, and — in the case of the outer cocoons propagating through an undisturbed IGM — the power-law index $\beta = 1.5$."
+ Such a density profile is expected to be a good approximation of the hot gaseous medium in groups of galaxies hosting luminous radio sources 1997)., Such a density profile is expected to be a good approximation of the hot gaseous medium in groups of galaxies hosting luminous radio sources \citep{KDA97}.
+. When modeling the inner lobes. which propagate within a uniform plasma of the ouler cocoons. we set instead ο=0 or 0.1 (as determined by a quality of the fits).," When modeling the inner lobes, which propagate within a uniform plasma of the outer cocoons, we set instead $\beta = 0$ or 0.1 (as determined by a quality of the fits)."
+ Note that (he assumption regarding a uniform distribution of the matter within the lobes is justilied bv the fact that the sound speed thereby has to be in general hish when compared to the dvnamical timescale involved (seeKaiserοἱal.2000)., Note that the assumption regarding a uniform distribution of the matter within the lobes is justified by the fact that the sound speed thereby has to be in general high when compared to the dynamical timescale involved \citep[see][]{kai00}.
+". Finally. we assumed the fraction of ihe non-radiating particles ""=0 for the inner lobes. and //=10 for the outer lobes. in order to account for the expected entrainment processes enriching (he evolved cocoons with ihe matter."," Finally, we assumed the fraction of the non-radiating particles $k^{\prime} = 0$ for the inner lobes, and $k^{\prime} = 10$ for the outer lobes, in order to account for the expected entrainment processes enriching the evolved cocoons with the matter."
+ The model lits return the following free parameters: the total jet kinetic power Qj. the densitv of the central plateau in the distribution of the ambient medium py. the injection," The model fits return the following free parameters: the total jet kinetic power $Q_{\rm j}$, the density of the central plateau in the distribution of the ambient medium $\rho_{0}$ , the injection"
+"As discussed in Sect. 4.1,,","As discussed in Sect. \ref{sub_basic_effects},"
+ the principle effects of clouds on the reflection spectra are the increases in both the reflected light and the related depths of absorption bands., the principle effects of clouds on the reflection spectra are the increases in both the reflected light and the related depths of absorption bands.
+" However, owing to the different cloud covers needed to obtain mean Earth surface temperature conditions the strength of these effects is different for each central star."," However, owing to the different cloud covers needed to obtain mean Earth surface temperature conditions the strength of these effects is different for each central star."
+" For example, the reflected radiation flux of the F-type star varies by about a factor of three between the minimum and maximum cloud coverage."," For example, the reflected radiation flux of the F-type star varies by about a factor of three between the minimum and maximum cloud coverage."
+" However, the reflection of light from the planet around the M-type star varies only slightly between its minimum and maximum cloud covers."," However, the reflection of light from the planet around the M-type star varies only slightly between its minimum and maximum cloud covers."
+ This is caused by the different cloud cover combinations required for the F star case and the M-type star to reproduce the global mean Earth surface temperature., This is caused by the different cloud cover combinations required for the F star case and the M-type star to reproduce the global mean Earth surface temperature.
+" For the M-type star, a large minimum coverage of low-level clouds is needed to cool the surface, whereas the required minimum high-level cloud cover is quite small in the F star case."," For the M-type star, a large minimum coverage of low-level clouds is needed to cool the surface, whereas the required minimum high-level cloud cover is quite small in the F star case."
+" As already mentioned, a wavelength-dependent albedo allows the detection of molecular absorption bands almost independently of the spectral energy distribution of the incident stellar radiation."," As already mentioned, a wavelength-dependent albedo allows the detection of molecular absorption bands almost independently of the spectral energy distribution of the incident stellar radiation."
+ The deep and broad bands of H5O are again visible in each case (Fig. 7))., The deep and broad bands of $\mathrm H_2 \mathrm O$ are again visible in each case (Fig. \ref{288_ratio}) ).
+" However, the Os band is not directly detectable in case of the smallest amount of clouds for the F high-level cloud cover) and G low-level cloud cover) stars."," However, the $\mathrm O_3$ band is not directly detectable in case of the smallest amount of clouds for the F high-level cloud cover) and G low-level cloud cover) stars."
+" Owing to the large amount of low-level clouds needed to reproduce the mean Earth surface temperature observed, the absorption bands of the biomarkers Os and O» can be seen for the M-type and K-type stars for every considered combination of cloud covers."," Owing to the large amount of low-level clouds needed to reproduce the mean Earth surface temperature observed, the absorption bands of the biomarkers $\mathrm O_3$ and $\mathrm O_2$ can be seen for the M-type and K-type stars for every considered combination of cloud covers."
+ The second axis on the right hand side of Fig., The second axis on the right hand side of Fig.
+ 7 shows in addition the contrast C; between the reflected planetary and incident stellar radiation fluxes., \ref{288_ratio} shows in addition the contrast $C_\lambda$ between the reflected planetary and incident stellar radiation fluxes.
+ The ratio between the radiation fluxes of the planet and the central star is quite small because of the much higher stellar luminosities compared to the planetary reflection spectra., The ratio between the radiation fluxes of the planet and the central star is quite small because of the much higher stellar luminosities compared to the planetary reflection spectra.
+" In the clear sky cases, the contrast varies between 1071? for the F-type star and 107? for the M-type star because the luminosity of the F-type star is much higher than that of the M-type star."," In the clear sky cases, the contrast varies between $10^{-10}$ for the F-type star and $10^{-8}$ for the M-type star because the luminosity of the F-type star is much higher than that of the M-type star."
+ The Earth-like planet around the M star is in addition located much closer to its host star as discussed in Paper I. Clouds do increase the contrast by about one order of magnitude owing to the already discussed enhancing scattering effects., The Earth-like planet around the M star is in addition located much closer to its host star as discussed in Paper I. Clouds do increase the contrast by about one order of magnitude owing to the already discussed enhancing scattering effects.
+" However, the resulting higher contrast values (10? type star - 1077 M-type star) are still very small."," However, the resulting higher contrast values $10^{-9}$ F-type star - $10^{-7}$ M-type star) are still very small."
+" Therefore, it is unlikely that these reflection spectra could be measured by the current generation of telescopes, even for the extremely high contrasts produced by clouds."," Therefore, it is unlikely that these reflection spectra could be measured by the current generation of telescopes, even for the extremely high contrasts produced by clouds."
+ We have combined a ID radiative-convective atmospheric model with our previously developed parametrised cloud description (Paper I) to study the influence of low and high-level clouds on the reflection spectra and spectral albedos of Earth-like extrasolar planets orbiting different types of central stars., We have combined a 1D radiative-convective atmospheric model with our previously developed parametrised cloud description (Paper I) to study the influence of low and high-level clouds on the reflection spectra and spectral albedos of Earth-like extrasolar planets orbiting different types of central stars.
+ In view of habitability NUV-NIR reflection spectra are of particular importance for the (possible) detection of the most important potential biomarker molecules Oy and O; (Fraunhofer A-band at 0.76um and the Chappuis band at 0.55—0.70 um)., In view of habitability NUV-NIR reflection spectra are of particular importance for the (possible) detection of the most important potential biomarker molecules $\mathrm O_2$ and $\mathrm O_3$ (Fraunhofer A-band at $0.76 \ \mathrm{\mu m}$ and the Chappuis band at $0.55-0.70 \ \mathrm{\mu m}$ ).
+ The detectability of molecular features in this wavelength regime is strongly supported by the presence of clouds., The detectability of molecular features in this wavelength regime is strongly supported by the presence of clouds.
+" Their high scattering efficiency substantially increases the reflected light and the related depths of the absorption bands, in comparison to the respective clear sky conditions where only the broad bands of H5O are visible."," Their high scattering efficiency substantially increases the reflected light and the related depths of the absorption bands, in comparison to the respective clear sky conditions where only the broad bands of $\mathrm H_2 \mathrm O$ are visible."
+ Low-level clouds have thereby a larger impact on the spectra than the high-level clouds because of their much higher scattering optical depth., Low-level clouds have thereby a larger impact on the spectra than the high-level clouds because of their much higher scattering optical depth.
+43.2). and NGC 1068 (logLx= 43.7).,"), and NGC 1068 $\log \LX = 43.7$ )."
+ A statistical analysis of the correlation. for our sample shows Spearman ranks PSpeannan—0.75...—0.9 with significance levels in the range of 0.2—5x1077., A statistical analysis of the correlation for our sample shows Spearman ranks $\rho_{\rm Spearman} = -0.75\ldots-0.9$ with significance levels in the range of $0.2-5\times10^{-2}$.
+ Despite the limited number of objects and luminosity coverage. correlations evolving out of the data can be confidently established.," Despite the limited number of objects and luminosity coverage, correlations evolving out of the data can be confidently established."
+ For BLBE and oNLBE studies. higher significance is achieved by averaging spectra of different objects within a narrow luminosity range.," For BLBE and oNLBE studies, higher significance is achieved by averaging spectra of different objects within a narrow luminosity range."
+ This overcomes peculiarities of individual objects (e.g.Croometal.2002)., This overcomes peculiarities of individual objects \citep[e.g.][]{Cro02}.
+. We aim for observations of a larger sample to make similar studies for the presented INLBE., We aim for observations of a larger sample to make similar studies for the presented iNLBE.
+" It is important to note that no correlation of W, with AGN distance. Dj. is present in our sample (Pspeannan<9.25. significance >0.6 for all lines)."," It is important to note that no correlation of $\EW$ with AGN distance, $D_L$ , is present in our sample $\rho_{\rm Spearman}<0.25$, significance $>0.6$ for all lines)."
+ As of yet. only one other study mentions a possible Baldwin effect in the nid-ifrared.," As of yet, only one other study mentions a possible Baldwin effect in the mid-infrared."
+ In an AAS abstract. Keremedjiev&Hao(2006) preserted the detection of a Baldwin effect for the [Siv] line in AGN data obtained by the satellite.," In an AAS abstract, \citet{Ker06} presented the detection of a Baldwin effect for the ] line in AGN data obtained by the satellite."
+ They also note indications of a Baldwin effect for n]. admitting that their study is suffering from the low spatial resolution of the data that they used.," They also note indications of a Baldwin effect for ], admitting that their study is suffering from the low spatial resolution of the data that they used."
+ In the available abstract. nothing is mentioned about a slope or the scatter of the anti-correlation.," In the available abstract, nothing is mentioned about a slope or the scatter of the anti-correlation."
+ Here. we demonstrate that the Baldwin effect of both the iv] and [Neu] line. and in addition the ut] line. is quite significant — and strong — when using high spatial resolution data. even with a small object sample.," Here, we demonstrate that the Baldwin effect of both the ] and ] line, and in addition the ] line, is quite significant – and strong – when using high spatial resolution data, even with a small object sample."
+ Thus. high spatial resolutio| appears to be crucial.," Thus, high spatial resolution appears to be crucial."
+ For the three type | AGN in our sample. black hole masses. Mpy. have been estimated based on reverberation mapping data.," For the three type 1 AGN in our sample, black hole masses, $\MBH$, have been estimated based on reverberation mapping data."
+ We adopt black hole masses for NGC 3227 (logMpu(Mq)=7.63+ 0.31) and NGC 3783 7.47+ 0.08) from Onkenetal.(2004).. and for NGC 4593 from Denneyetal.(2006).," We adopt black hole masses for NGC 3227 $\log\MBH(\Msun) = 7.63 \pm 0.31$ ) and NGC 3783 $\log\MBH = 7.47 \pm 0.08$ ) from \citet{Onk04}, and for NGC 4593 from \citet{Den06}."
+. In addition. the black hole mass of the Seyfert 2 galasy NGC 1068 could be determined by MASER cloud kinematics as logμμ7.0 (Greenhilletal.1996).," In addition, the black hole mass of the Seyfert 2 galaxy NGC 1068 could be determined by MASER cloud kinematics as $\log\MBH \approx 7.0$ \citep{Gre96}."
+. For these four AGN. we analyzed the dependence of the equivalent widths of the [Arm] liπο on Mpypg.," For these four AGN, we analyzed the dependence of the equivalent widths of the ] line on $\MBH$."
+ A statistical test shows no evident correlation for our limited sample of objects and the small coverage of black hole Masses (Pspeaman=0.40. significance 0.6. for a nominal fit W(Armpe/Duc ).," A statistical test shows no evident correlation for our limited sample of objects and the small coverage of black hole masses $\rho_{\rm Spearman}=0.40$, significance $0.6$, for a nominal fit $\EW$ $\propto \MBH^{0.5\pm0.7}$ )."
+ Subject to the limitations. this indicates a fundamental difference of the INLBE as compared to the BLBE for which such an anti-correlation has been found (Warneretal.2004).," Subject to the limitations, this indicates a fundamental difference of the iNLBE as compared to the BLBE for which such an anti-correlation has been found \citep{War04}."
+ By using bolometric correction to the X-ray luminosities listed in Table | (Marcontetal. 2004).. itis possible to estimate the Eddington ratio. γω. of four AGN with known Mpy.," By using bolometric correction to the X-ray luminosities listed in Table \ref{Tab:EW} \citep{Mar04}, , it is possible to estimate the Eddington ratio, $\ledd$, of four AGN with known $\MBH$."
+" A statistical test for a correlation between W([ Arm) and {τμ reveals a Spearman rank Ουμας=0.5 with a significance of 0.2 for the relation W,(Arumtpoc{πμ .", A statistical test for a correlation between $\EW($ $)$ and $\ledd$ reveals a Spearman rank $\rho_{\rm Spearman} = -0.8$ with a significance of 0.2 for the relation $\EW$ $\propto \ledd^{-0.40\pm0.17}$ .
+" Thus. our limited sample doesn't allow us to firmly establish a correlation. but a negative dependence of W, on {μμ is indicated."," Thus, our limited sample doesn't allow us to firmly establish a correlation, but a negative dependence of $\EW$ on $\ledd$ is indicated."
+ We excluded Circinus from all analysis of the correlations., We excluded Circinus from all analysis of the correlations.
+ As can be seen in Figs., As can be seen in Figs.
+" | 2.. the W,’s of all three mid-infrared lines deviate significantly from the nominal fit to the other AGN."," \ref{Fig:LineComb} \ref{Fig:NeII}, the $\EW$ 's of all three mid-infrared lines deviate significantly from the nominal fit to the other AGN."
+ Although it could be a resolution effect (Circinus is the nearest AGN in the sample). NGC 1068 has about the same spatial resolutiol when scaled for the different luminosities (scaling r«L! *).," Although it could be a resolution effect (Circinus is the nearest AGN in the sample), NGC 1068 has about the same spatial resolution when scaled for the different luminosities (scaling $r\propto L^{1/2}$ )."
+ Toreover. we did not detect extended emission in the lines down to ~5% of the peak flux.," Moreover, we did not detect extended emission in the lines down to $\sim$ of the peak flux."
+ This gives us some confiderce toreject that a significant amount of line flux is over-resolved and lost in the noise., This gives us some confidence toreject that a significant amount of line flux is over-resolved and lost in the noise.
+Circinus is khown for strong dust extinction towards the nucleusby dust lanes in both the host and our own Galaxy.,Circinus is known for strong dust extinction towards the nucleusby dust lanes in both the host and our own Galaxy.
+ In, In
+also Fig 4)) and putting both timescales equal. this yields the transition frequency v. of 10.1 GHz.,"also Fig \ref{fig4}) ) and putting both timescales equal, this yields the transition frequency $\nu_c$ of 10.1 GHz."
+ Below this frequency PSR B0329+54 will be in the strong scintillation regime. above that frequency. the characteristic of the scintillation should change to weak seintillations.," Below this frequency PSR B0329+54 will be in the strong scintillation regime, above that frequency, the characteristic of the scintillation should change to weak scintillations."
+ Observations of the interstellar scintillations and especially the frequency dependence of the measured parameters can be usec to estimate the properties of the interstellar turbulence., Observations of the interstellar scintillations and especially the frequency dependence of the measured parameters can be used to estimate the properties of the interstellar turbulence.
+" By analogy to the neutral gas theory. the density fluctuations Ἡ the ISM can be deseribed by a power-law spectrum as Ps,=C;4 q?. where C7 is the mean turbulence of electron density along the line of sight. q=27/5 is the wavenumber associatec with the spatial scale of turbulence s. and the spectral index P. which is in range of 3«f<$5."," By analogy to the neutral gas theory, the density fluctuations in the ISM can be described by a power-law spectrum as $P_{\rm 3n}=C^{2}_{n}~q^{-\beta}$ , where $C^2_n$ is the mean turbulence of electron density along the line of sight, $q=2\pi/s$ is the wavenumber associated with the spatial scale of turbulence $s$, and the spectral index $\beta$, which is in range of $3<\beta<5$."
+ The well knowt Kolmogorov theory developed for neutral gas turbulence yields the spectral index of B=11/3. but it was shown by several authors (see for example Gupta et al. 1993))," The well known Kolmogorov theory developed for neutral gas turbulence yields the spectral index of $\beta = 11/3$, but it was shown by several authors (see for example Gupta et al. \cite{gupt}) )"
+ that there are discrepancies between Kolmogorov theoretical predictions and actual measurements of scintillation parameters for some pulsars., that there are discrepancies between Kolmogorov theoretical predictions and actual measurements of scintillation parameters for some pulsars.
+ Few models of the interstellar turbulence with different spectral slopes ϱ were developed. yielding different sets of predictions of the frequency dependence of ISS parameters (see Romani et al. 1986:;," Few models of the interstellar turbulence with different spectral slopes $\beta$ were developed, yielding different sets of predictions of the frequency dependence of ISS parameters (see Romani et al. \cite{romani};"
+ Bhat et al. 1999)., Bhat et al. \cite{bhat}) ).
+ Table 4. summarizes these predictions for the three most commonly used spectral slopes ofB= 11/3 (Kolmogorov spectrum). 4 and 4.3 for the four scintillation parameters we were able to measure. or derive from our observations of PSR BO329+54.," Table \ref{indx_table} summarizes these predictions for the three most commonly used spectral slopes of $\beta = $ 11/3 (Kolmogorov spectrum), 4 and 4.3 for the four scintillation parameters we were able to measure, or derive from our observations of PSR B0329+54."
+ Our estimations of the actual spectral slopes of these parameters. based on all available data (see section 3.2. and Figure 4)) are also in the table.," Our estimations of the actual spectral slopes of these parameters, based on all available data (see section \ref{scint_params} and Figure \ref{fig4}) ) are also in the table."
+ Apparently none of the currently. models for the electron turbulence agrees precisely with our estimations. and the model predictions with8=4 are the closest to the slopes we obtained from our fits.," Apparently none of the currently models for the electron turbulence agrees precisely with our estimations, and the model predictions with $\beta=4$ are the closest to the slopes we obtained from our fits."
+ It has to be noted that in case of three of the four parameters (with the sipiss being the only exception). these discrepancies cannot be attributed to poor quality of frequency slopes fits for these parameters. especially that our measurements. made at the frequency of 4.8 GHz. significantly widened the frequency range. putting strong constraints on the fits.," It has to be noted that in case of three of the four parameters (with the $m_{\rm RISS}$ being the only exception), these discrepancies cannot be attributed to poor quality of frequency slopes fits for these parameters, especially that our measurements, made at the frequency of 4.8 GHz, significantly widened the frequency range, putting strong constraints on the fits."
+ As we mentioned above. in a strong scintillation regime the flux density variations happen at two distinctive timescales - diffractive and refractive.," As we mentioned above, in a strong scintillation regime the flux density variations happen at two distinctive timescales - diffractive and refractive."
+ The apparent variability of the pulsar signal can be understood as the result of the observer crossing the diffraction pattern. which ts introduced by the ISM to the pulsar signal wavefront.," The apparent variability of the pulsar signal can be understood as the result of the observer crossing the diffraction pattern, which is introduced by the ISM to the pulsar signal wavefront."
+" In a thin screen model. those two timescales correspond to two spatial scales of the diffraction pattern: the diffractive scale sj. and refractive scale s,. which in turn can be bound to the concepts of scattering angle and scattering disk that is commonly used in the scintillation theory (see Rickkett. 1990)."," In a thin screen model, those two timescales correspond to two spatial scales of the diffraction pattern: the diffractive scale $s_d$, and refractive scale $s_r$, which in turn can be bound to the concepts of scattering angle and scattering disk that is commonly used in the scintillation theory (see Rickkett, \cite{rick}) )."
+ Following Gupta et al. (1994)).," Following Gupta et al. \cite{gup94}) ),"
+ we can find the diffractive scale by the means of estimating the scintillation velocity Viss.," we can find the diffractive scale by the means of estimating the scintillation velocity $V_{\rm ISS}$ ,"
+Al: ~3 the galaxy. contribution has been recently evaluated by SPA who adopted their estimate of the 1500 oealaxy luminosity function and their derived escape fraction of65.,At $z\sim 3$ the galaxy contribution has been recently evaluated by SPA who adopted their estimate of the 1500 galaxy luminosity function and their derived escape fraction of.
+".. Their resultinee UVB intensity [rom the egalaxy population alone is ./=1.2x102l erese i ? Fay to which is already. larger than the value estimated bv the ""proximity. effect in the Lyman-a forest of QSO absorption spectra. J=5—7x1077 eres !H 7? H | (see e.g. Giallongo et al."," Their resulting UVB intensity from the galaxy population alone is $J=1.2\times 10^{-21}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ which is already larger than the value estimated by the ""proximity effect"" in the $\alpha$ forest of QSO absorption spectra, $J=5-7\times 10^{-22}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ (see e.g. Giallongo et al."
+ 1996: Scott. et al., 1996; Scott et al.
+ 2000)., 2000).
+ The addition of the contribution by the QSO population would increase (his discrepancy., The addition of the contribution by the QSO population would increase this discrepancy.
+ To alleviate the problem thev argued about a possible overestimate of (he QSO contribution to the UVB., To alleviate the problem they argued about a possible overestimate of the QSO contribution to the UVB.
+ Dianchi et al. (, Bianchi et al. (
+"2001) used the observed UV star formation history and three different values for the intrinsic fi300/foun ralio. namely 5.3. 21. 42 corresponding to. fi4,4.=5T.1H.1.. respectively. (following the definition of f,,,:. adopted in the present paper) to compare the contribution of the two populations to the UVB and the effects on the evolution of the Lyman-a forest.","2001) used the observed UV star formation history and three different values for the intrinsic $f_{1500} /f_{900}$ ratio, namely 5.3, 21, 42 corresponding to $f_{rel,esc} =
+57, 14, 7$, respectively, (following the definition of $f_{rel,esc}$ adopted in the present paper) to compare the contribution of the two populations to the UVB and the effects on the evolution of the $\alpha$ forest."
+ It is found that. to avoid an overprediction of the total (QSO0-- galaxies) UVB respect to the value derived from the proximity effect. the relative escape fraction should be of the order of or smaller (han1556.," It is found that, to avoid an overprediction of the total $+$ galaxies) UVB respect to the value derived from the ""proximity effect"", the relative escape fraction should be of the order of or smaller than."
+. In this respect our upper limit supports a scenario where the total UVB is consistent with the “proximity elfect al z~3 with a galaxy contribution of JX6x107 eres tem 7 PF J|. iie; no more than a [actor 2.5 higher than the QSO contribution.," In this respect our upper limit supports a scenario where the total UVB is consistent with the ""proximity effect"" at $z\sim 3$ with a galaxy contribution of $J\lesssim 6\times 10^{-22}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$, i.e. no more than a factor 2.5 higher than the QSO contribution."
+ An UVB produced by this mixed population could account for the metal enrichment of the IGM and for the SilV/CIV metal-line ratios observed in hieh resolution QSO absorption spectra (see e.g. Giroux Shull 1991)., An UVB produced by this mixed population could account for the metal enrichment of the IGM and for the SiIV/CIV metal-line ratios observed in high resolution QSO absorption spectra (see e.g. Giroux Shull 1997).
+ In summary. many pieces of evidence. from the limits derived at lower and intermediate redshilts to the upper limit of {εως£16% presented in the present paper. point toward a lower ionizing escape fraction.," In summary, many pieces of evidence, from the limits derived at lower and intermediate redshifts to the upper limit of $f_{rel,esc} \lesssim 16$ presented in the present paper, point toward a lower ionizing escape fraction."
+ We emphasize that our result. however robust. is based on the observations of two z~3 galaxies only.," We emphasize that our result, however robust, is based on the observations of two $z\sim 3$ galaxies only."
+ A large statistical sample is needed to confirm iis implication on the UV background evolution., A large statistical sample is needed to confirm its implication on the UV background evolution.
+ This is within the grasp of the current instrumentation al verv large telescopes., This is within the grasp of the current instrumentation at very large telescopes.
+ We are grateful to C. C. Steidel. M. Pettind and collaborators for providing the finding charts of the two galaxies.," We are grateful to C. C. Steidel, M. Pettini and collaborators for providing the finding charts of the two galaxies."
+It is important that values derived from the results of the model fitter include several models. in order to obtain a reasonable estimate of whether these results are constrained and unique with the data used.,"It is important that values derived from the results of the model fitter include several models, in order to obtain a reasonable estimate of whether these results are constrained and unique with the data used."
+ Therefore the number of fits within a Ay*=2 of the best fit was calculated (weusethesamedefinitionof y-as2007)., Therefore the number of fits within a $\Delta\chi^{2}$ =2 of the best fit was calculated \citep[we use the same definition of $\chi^{2}$ .
+ If there were less than 10 fits included within Ves + Ay? then Ay? was increased by integer steps until this was no longer the case., If there were less than 10 fits included within $\chi_{best}^{2}$ + $\Delta\chi^{2}$ then $\Delta\chi^{2}$ was increased by integer steps until this was no longer the case.
+ This method was used rather than simply finding the best 10 fits so that the spread 1 fits was not artificially limited if many models are good fits to the data with a range of values for distance. bolometric flux and Ay.," This method was used rather than simply finding the best 10 fits so that the spread in fits was not artificially limited if many models are good fits to the data with a range of values for distance, bolometric flux and $A_{V}$."
+ A minimum of 10 fits were required in order to esure that one or two fits did not strongly bias the results., A minimum of 10 fits were required in order to ensure that one or two fits did not strongly bias the results.
+ The model bolometric flux for the source was obtained by using the model luminosity and model distance. derived for each individual fit to calculate the bolometric flux for each fit (f;)., The model bolometric flux for the source was obtained by using the model luminosity and model distance derived for each individual fit to calculate the bolometric flux for each fit $f_{i}$ ).
+ The luminosity returned by the model fitter is the combination of that for the central source. accretion onto the central source and accretion in the disk. which is then corrected for foreground extinction consistent with the SED. and is therefore independent of viewing angle unlike the SED itself.," The luminosity returned by the model fitter is the combination of that for the central source, accretion onto the central source and accretion in the disk, which is then corrected for foreground extinction consistent with the SED, and is therefore independent of viewing angle unlike the SED itself."
+ Next the weighted mean and weighted standard deviation were calculated of all fits within Ves Αγ; to give the flux and flux uncertainty using: where N' is the number of non-zero weights. and the weights are calculated from the y as: The y values are used to weight the fits so that the best model fitted to the data has more impact on the results than poorer fits. but that the results are not heavily skewed away from the mean.," Next the weighted mean and weighted standard deviation were calculated of all fits within $\chi_{best}^{2}$ + $\Delta\chi^{2}$ to give the flux and flux uncertainty using: where $N^{'}$ is the number of non-zero weights, and the weights are calculated from the $\chi^{2}$ as: The $\chi^{2}$ values are used to weight the fits so that the best model fitted to the data has more impact on the results than poorer fits, but that the results are not heavily skewed away from the mean."
+ This method ts also used to calculate the mean model distance. the mean Ay. and mean filter flux ratios given the results for individual model fits to the data.," This method is also used to calculate the mean model distance, the mean $A_{V}$, and mean filter flux ratios given the results for individual model fits to the data."
+ A distance range tis required by the model fitter within which to fit the observed data., A distance range is required by the model fitter within which to fit the observed data.
+ Therefore the test described below was undertaken in order to explore the distance dependence of the bolometric flux obtained from the model fitter. and thus what effect an incorrect distance has on themodel derived flux.," Therefore the test described below was undertaken in order to explore the distance dependence of the bolometric flux obtained from the model fitter, and thus what effect an incorrect distance has on themodel derived flux."
+ Fits were performed to the data for GO34.7569+00.0247 for model distances ¢ + 6d kpe. with d varying between 2 kpc and 13 kpe in I kpe integer steps.," Fits were performed to the data for G034.7569+00.0247 for model distances $d$ $\pm$ $\delta$$d$ kpc, with $d$ varying between 2 kpc and 13 kpc in 1 kpc integer steps."
+ Example SED fits for 2.0 + 1.0 kpe and 5.0 + 1.0 kpe (thekinematicdistanceG034.7569+00.0247is5.0+0.7kpe.Urquhartetal..20082) are shown in Figure | (left and right respectively).," Example SED fits for 2.0 $\pm$ 1.0 kpc and 5.0 $\pm$ 1.0 kpc \citep[the kinematic distance for G034.7569+00.0247 is 5.0~$\pm$~0.7~kpc,][]{Urquhart2008a} are shown in Figure \ref{F:tests_distance_seds} (left and right respectively)."
+ For all runs of the model fitter. od is set to 1.0 kpe as this ts the order of the typical error in the kinematic distance for RMS sources.," For all runs of the model fitter, $\delta$$d$ is set to 1.0 kpc as this is the order of the typical error in the kinematic distance for RMS sources."
+ The model fitter divides the user-specified distance range logarithmically in order to obtain model distances after which the model SED fluxes are convolved with common filter bands and interpolated to the apertures used to obtain the data., The model fitter divides the user-specified distance range logarithmically in order to obtain model distances after which the model SED fluxes are convolved with common filter bands and interpolated to the apertures used to obtain the data.
+ The logarithmie binning was set to 0.01 for this test. while 0.025 was used for the other tests and the results presented in refS:results..," The logarithmic binning was set to 0.01 for this test, while 0.025 was used for the other tests and the results presented in \\ref{S:results}."
+ The results were checked for a few sources at the lower model distance binning step and the differences were found to be minimal., The results were checked for a few sources at the lower model distance binning step and the differences were found to be minimal.
+ At small model distances the apertures may not encompass the whole model but only the inner region. the calculation of which can require large amounts of computer memory.," At small model distances the apertures may not encompass the whole model but only the inner region, the calculation of which can require large amounts of computer memory."
+ Therefore. where the distance to the source was less than 2.4 kpc. the distance given to the fitter for the main results Was reset to this minimum limit.," Therefore, where the distance to the source was less than 2.4 kpc, the distance given to the fitter for the main results was reset to this minimum limit."
+ The results of the fits where d is varied for the data of G034.7569-00.0247 are shown in Figure 2.. in terms of bolometric flux (left) and Ay (right) vs. the mean model distance as calculated using the weighted mean and weighted standard deviation of the model distances for the individual fits within Ay? of TAM," The results of the fits where $d$ is varied for the data of G034.7569+00.0247 are shown in Figure \ref{F:tests_distance_allruns}, in terms of bolometric flux (left) and $A_{V}$ (right) vs. the mean model distance as calculated using the weighted mean and weighted standard deviation of the model distances for the individual fits within $\Delta$$\chi^{2}$ of $\chi^{2}_{best}$."
+ The A 15 relatively constant with distance for these fits (see Figure 2)). suggesting it is constrained by the data.," The $A_{V}$ is relatively constant with distance for these fits (see Figure \ref{F:tests_distance_allruns}) ), suggesting it is constrained by the data."
+ The increase in the bolometric flux and error at small distances is due to the source being so close that the Ay is insufficient to stop emission in the visual being observed (see Figure 1))., The increase in the bolometric flux and error at small distances is due to the source being so close that the $A_{V}$ is insufficient to stop emission in the visual being observed (see Figure \ref{F:tests_distance_seds}) ).
+ The model fitter allows this as the near-IR fluxes are upper limits. and so are treated as weaker constraints on the SED than the detections at longer wavelengths.," The model fitter allows this as the near-IR fluxes are upper limits, and so are treated as weaker constraints on the SED than the detections at longer wavelengths."
+ Overall. though the bolometric flux appears to decrease with increasing distance. the change 1s of similar order or smaller than the uncertainties in the fluxes. and so the results can be considered distance independent.," Overall, though the bolometric flux appears to decrease with increasing distance, the change is of similar order or smaller than the uncertainties in the fluxes, and so the results can be considered distance independent."
+ The visual extinction. CA). between Earth and the edge of the circumstellar material considered by the model is an important free parameter in fitting model SEDs to observed data., The visual extinction $A_{V}$ ) between Earth and the edge of the circumstellar material considered by the model is an important free parameter in fitting model SEDs to observed data.
+ Restricting this value too much could result in poor fits. particularly between the near/mid-IR and far-IR/sub-mm parts of the SED. while allowing too much freedom could lead to unphysical solutions.," Restricting this value too much could result in poor fits, particularly between the near/mid-IR and far-IR/sub-mm parts of the SED, while allowing too much freedom could lead to unphysical solutions."
+ The extinetion law model assumed was calculated by Robitailleetal.(2007) using the method of Kimetal.(1994) for a typical galactic ISM. taking into account the results of Indebetouwetal.(2005) with regards to mid-IR extinction.," The extinction law model assumed was calculated by \citet{Robitaille2007a} using the method of \citet{Kim1994} for a typical galactic ISM, taking into account the results of \citet{Indebetouw2005} with regards to mid-IR extinction."
+ We limit the Ay to be within the range 0 to 50. in line with typical values for regions (e.g.Hansonetal..2002).. though as shown in Figure 3 the majority of sources have fitted values below this.," We limit the $A_{V}$ to be within the range 0 to 50, in line with typical values for regions \citep[e.g.][]{Hanson2002}, though as shown in Figure \ref{F:tests_extinction} the majority of sources have fitted values below this."
+ SED model fits for GO34.7569+00.0247 including only subsets of the available data were undertaken in order to explore the effect of missing data on the bolometric flux results. examples of which. along with the fit using all data are shown in Figure 4..," SED model fits for G034.7569+00.0247 including only subsets of the available data were undertaken in order to explore the effect of missing data on the bolometric flux results, examples of which, along with the fit using all data are shown in Figure \ref{F:tests_data_seds}. ."
+ Comparing the top plot. which includes all data. with the middle plot where the MIPSGAL data has been omitted. the effect of using the model fitter for sources without far-IR data," Comparing the top plot, which includes all data, with the middle plot where the MIPSGAL data has been omitted, the effect of using the model fitter for sources without far-IR data"
+a coupling between Luectuations on clillerent time-scales. ruling out models where the variability arises [from strictly independent. [lares or active regions(Utthev.MIHIardy&Vaughan 2005).,"a coupling between fluctuations on different time-scales, ruling out models where the variability arises from strictly independent flares or active regions\citep{UttleyNonlin}."
+. Fluctuating accretion mocoels. on the other hand. are consistent with this relation: the Duetuations couple together as they propagate down to the centre. so low frequeney Buctuations produced. at large radii modulate higher. frequency. Huctuations. produced. further in.," Fluctuating accretion models, on the other hand, are consistent with this relation: the fluctuations couple together as they propagate down to the centre, so low frequency fluctuations produced at large radii modulate higher frequency fluctuations produced further in."
+ A linear rms-Dux. relation is also observed in neutron star NRB svstenis.. where the emission. mechanism is most likely dillerent. to the one operating in.black-hole systems (Uttley2004:Uttles&AlHardy2001)..," A linear rms-flux relation is also observed in neutron star XRB systems, where the emission mechanism is most likely different to the one operating inblack-hole systems \citep{Uttley_sax,Uttley}."
+. This result provides aclelitional evidence for Huctuating-accretion scenarios because it suggests that the variability originates in the aceretion Dow and not just locally via the mechanism producing the emission (Uttley2004)., This result provides additional evidence for fluctuating-accretion scenarios because it suggests that the variability originates in the accretion flow and not just locally via the mechanism producing the emission \citep{Uttley_sax}.
+. 3esicles explaining the broad range of variability time-scales and the rms-Dux. relation. the Iluctuating accretion model can also explain the spectral-timing properties of the variability. as noted by Ποιονetal.(2001).," Besides explaining the broad range of variability time-scales and the rms-flux relation, the fluctuating accretion model can also explain the spectral-timing properties of the variability, as noted by \citet{Kotov}."
+.. For example. it has long been known (e.g.Alivamoto&Witamoto1980:Nowaketal.1999a). that N-rav. variations in black hole X- binaries (BIINRBs) often show hard lags. Le. a delay in hard. X-ray variations with respect to soft. N-ravs. which is larger for variations on longer time-scales. and at higher energies.," For example, it has long been known \citep[e.g.][]{Miyamoto,NowakGX339} that X-ray variations in black hole X-ray binaries (BHXRBs) often show hard lags, i.e. a delay in hard X-ray variations with respect to soft X-rays, which is larger for variations on longer time-scales, and at higher energies."
+ The magnitudes of the lags are typically of order one per cent of the variability time-scale., The magnitudes of the lags are typically of order one per cent of the variability time-scale.
+ Similar time-scale-dependent. hard. lags have recently been discovered in AGN. albeit on much longer time-scales. commensurate with their higher black hole masses (c.g.Papaclakisetal.2001:Vaughanctal.2003:MLlardy.et 2004).," Similar time-scale-dependent, hard lags have recently been discovered in AGN, albeit on much longer time-scales, commensurate with their higher black hole masses \citep[e.g.][]{Papadakis,VaughanMCG,McHardy4051}."
+. In their analvtical extension to Lvubarskiis model. Ixotov.et.al.(2001) explain these lags by invoking a raclially extended A-ray emitting region with an encrey-cdependent profile.," In their analytical extension to Lyubarskii's model, \citet{Kotov}
+ explain these lags by invoking a radially extended X-ray emitting region with an energy-dependent profile."
+ Ln his scenario. the response of the emission to a fluctuation in the aceretion How is a function of the inward propagation ime-scale. anc hence radius of origin. of the Uuctuation. combined with the emissivity profile.," In this scenario, the response of the emission to a fluctuation in the accretion flow is a function of the inward propagation time-scale, and hence radius of origin, of the fluctuation, combined with the emissivity profile."
+ Llenee. if the emissivity oofile is more centrally concentrated: at. higher energies. iard-band. lags are produced such that the lags are larger or longer time-scale variations.," Hence, if the emissivity profile is more centrally concentrated at higher energies, hard-band lags are produced such that the lags are larger for longer time-scale variations."
+ Ixotov.etal.(2001). show hat the same basic picture can also explain the energy dependence of the PSD of BUINR and AGN. where there is relatively more high-frequency Bspower observed at higher energies than at lower energies (c.g.Nowakctal.1999b:ΑΕLardyctal. 2004).," \citet{Kotov} show that the same basic picture can also explain the energy dependence of the PSD of BHXRBs and AGN, where there is relatively more high-frequency power observed at higher energies than at lower energies \citep[e.g.][]{Nowak_lags,McHardy4051}."
+. This is because the emitting region acts as a low-pass filter and. as the emitting region of the soft band is more extended. the higher frequeney variations are filtered more stronely in the soft band.," This is because the emitting region acts as a low-pass filter and, as the emitting region of the soft band is more extended, the higher frequency variations are filtered more strongly in the soft band."
+ Due to its success in. explainiris many aspects of the X-ray variability data. the Ποιαing accretion mocel of Lyubarskii(1997)— and Ιουνeal.(2001) warrants a deeperinvestigation. which is he aim of this paper.," Due to its success in explaining many aspects of the X-ray variability data, the fluctuating accretion model of \citet{Lyub} and \citet{Kotov} warrants a deeperinvestigation, which is the aim of this paper."
+ As shown bv the brief. analytical treatment of Wotovetal. (2001).. in principle. the model produces time-scale- lags of approximately he same amplitude as observed in the data. and encrev-cledent PSDs.," As shown by the brief analytical treatment of \citet{Kotov}, in principle, the model produces time-scale-dependent lags of approximately the same amplitude as observed in the data, and energy-dependent PSDs."
+ Lowever it is not clear exactly how the lag spectrum (lag as a function of Fourier frequency) and energy. dependence of the PSD varies as a function of emissivity profile. or the parameters of the accretion Dow. (assuming the standard disc model of Shakura&Sunvaev1973. e.g. the viscosity parameter à or he scale-height of the How). or due to the cllects of racial damping of variations in the accretion How.," However it is not clear exactly how the lag spectrum (lag as a function of Fourier frequency) and energy dependence of the PSD varies as a function of emissivity profile, or the parameters of the accretion flow (assuming the standard disc model of \citealt{Shakura} e.g. the viscosity parameter $\alpha$ or the scale-height of the flow), or due to the effects of radial damping of variations in the accretion flow."
+ Furthermore. since Ixotovetal.(2001). analytically determined spectral-iming properties by making the simplifving assumption hat the perturbation introduced. into the accretion. [ow at each radius is a cdelta-function in time and radius. it is important to determine the effects. on. spectral-timing woperties of a more realistic mocel. where stochastic variations over a broader range of time-scales are introduced ab each radius.," Furthermore, since \citet{Kotov}
+ analytically determined spectral-timing properties by making the simplifying assumption that the perturbation introduced into the accretion flow at each radius is a delta-function in time and radius, it is important to determine the effects on spectral-timing properties of a more realistic model, where stochastic variations over a broader range of time-scales are introduced at each radius."
+ Finally. the Uuctuating accretion model should. in principle. produce light curves which are highly correlated. in different bands. (ie. spectrally coherent. sec Vaughan&Nowak(1997) ancl Appendix D of this paper). but this aspect of the model has not vet been studied.," Finally, the fluctuating accretion model should, in principle, produce light curves which are highly correlated in different bands (i.e. spectrally coherent, see \cite{Vaughan_coh} and Appendix B of this paper), but this aspect of the model has not yet been studied."
+ To study these various effects in more detail. in Section 2 we introduce a computational tov model for a [uctuating accretion [low. based on the work of Lyubarskii(1997) and Ixotovetal.(2001)... and explore the dependence of the PSD. time lags and coherence on the model parameters in Section 3..," To study these various effects in more detail, in Section \ref{model} we introduce a computational toy model for a fluctuating accretion flow, based on the work of \citet{Lyub} and \citet{Kotov}, and explore the dependence of the PSD, time lags and coherence on the model parameters in Section \ref{properties}."
+ We show how the model predictions fit X-ray data from ACN and the BUNRB €ve X-1 in Sections 4. and 5.., We show how the model predictions fit X-ray data from AGN and the BHXRB Cyg X-1 in Sections \ref{agn} and \ref{xrb}.
+ Specifically. we will concentrate on explaining the spectral-timing properties of the DIIXIUB Cyve N-1 in its high/soft state. and AGN which show similar variability. properties. since these states show rather simple Lff-type PSD shapes. without complex quasi-periocic oscillations (QPOs).," Specifically, we will concentrate on explaining the spectral-timing properties of the BHXRB Cyg X-1 in its high/soft state, and AGN which show similar variability properties, since these states show rather simple $1/f$ -type PSD shapes, without complex quasi-periodic oscillations (QPOs)."
+ We discuss the implications of our results in Section 6 and sunimarise our results in Section 7.., We discuss the implications of our results in Section \ref{discussion} and summarise our results in Section \ref{conclusions}.
+ The model is based on the scenario proposed by Lyubarskii (LOOT). where fluctuations propagate inward through the accretion Dow. modulating the emission of the inner regions.," The model is based on the scenario proposed by \citet{Lyub}, where fluctuations propagate inward through the accretion flow, modulating the emission of the inner regions."
+ The Uuetuations are. produced on. time-scales related: to he viscous time-scale at the radius of origin ancl are uncorrelatecl for cilferent radial scales., The fluctuations are produced on time-scales related to the viscous time-scale at the radius of origin and are uncorrelated for different radial scales.
+ Churazovctal.(2001) noted that. in this class of models. the accretion rate luctuations can be carried ellicientlv. and be produced up to ugh Fourier frequencies. bv a geometrically thick accretion low. identifying it with an extended corona over the thin disc.," \citet{Churazov} noted that, in this class of models, the accretion rate fluctuations can be carried efficiently, and be produced up to high Fourier frequencies, by a geometrically thick accretion flow, identifying it with an extended corona over the thin disc."
+ We use standard aceretion cise (Shakura&Sunvaev1973) considerations to relate the fluctuation time-scales ancl propagation speeds to the radial position in the clise. since the same relations are apjxicable to geometrically thick and optically thin accretion {ows. as well as the standard ecometrically thin disc.," We use standard accretion disc \citep{Shakura} considerations to relate the fluctuation time-scales and propagation speeds to the radial position in the disc, since the same relations are applicable to geometrically thick and optically thin accretion flows, as well as the standard geometrically thin disc."
+ Following Lyubarskii(1007)... we assume that cach independen annulus produces a pattern of Lluetuations m(r.£). having most of the variability power at the local viscous frequeney fos)=4πειταfox (e.gIxatoetal. 1998).. where (12/2) is the disc scale height to radius ratio. à is the viscosity parameter. the radial position r Is in units of gravtational radii Re=CALfc and thefrequency is given in ternis of ο...," Following \citet{Lyub}, we assume that each independent annulus produces a pattern of fluctuations $\dot m(r,t)$, having most of the variability power at the local viscous frequency $f_{\rm visc}(r)=r^{-3/2}(H/R)^2 \alpha/2\pi$ \citep[e.g][]{Kato}, , where $(H/R)$ is the disc scale height to radius ratio, $\alpha$ is the viscosity parameter, the radial position $r$ is in units of gravitational radii $R_{\rm g}=GM/c^2$ and thefrequency is given in terms of $c/R_{\rm g}$ ."
+ Phe local accretion rate at any racius in the disc is allowed to Ductuate around a value AL hl). às AG.{j=ALQo(01|mn. D). where mr.<< l," The local accretion rate at any radius in the disc is allowed to fluctuate around a value $\dot M_o(r,t)$ , as $\dot M(r,t)=\dot
+M_o(r,t)\times(1+\dot m (r,t))$ , where $\dot m(r,t) <<1$ ."
+o As this new value of the accretion rate propagates inward. it serves as the ALU.) for the," As this new value of the accretion rate propagates inward, it serves as the $\dot
+M_o(r,t)$ for the"
+Radiation observed from astrophysical systems like GRBs or AGN usually possesses a nonthermal emission spectrum.,Radiation observed from astrophysical systems like GRBs or AGN usually possesses a nonthermal emission spectrum.
+ This is believed to arise [rom particle acceleration in the vicinity of relativistie shocks or within the counterstreaming plasma itself., This is believed to arise from particle acceleration in the vicinity of relativistic shocks or within the counterstreaming plasma itself.
+ In recent PIC simulations it has been shown that most parücle acceleration occurs within the jet (?).. B," In recent PiC simulations it has been shown that most particle acceleration occurs within the jet \citep{2003ApJ...595..555N, 2005ApJ...622..927N, 2006ApJ...642.1267N, 2008AIPC.1085..589N, 2009arXiv0901.4058N, 2008ApJ...674..378C, 2008ApJ...673L..39S, 2008ApJ...675..586D, 2004ApJ...608L..13F, 2004ApJ...617L.107H, 2005ApJ...623L..89H, 2009ApJ...695L.189M, 2007ApJ...671.1877R, 2003ApJ...596L.121S} and is mostly caused by plasma instabilities like the Weibel \citep{1959PhRvL...2...83W} or twostream instability \citep{1958PhRvL...1....8B}."
+oth instabilities create current filaments with surrounding magnetic fields and are therefore a plausible source for parücle acceleration and the generation of observed long-lasting magnetic fields., Both instabilities create current filaments with surrounding magnetic fields and are therefore a plausible source for particle acceleration and the generation of observed long-lasting magnetic fields.
+" Particle acceleration can also occur along with shocks where first order Fermi acceleration (2). is assumed to be the relevant process, which was shown by kinetic simulations only recently (2).."," Particle acceleration can also occur along with shocks where first order Fermi acceleration \citep{1949PhRv...75.1169F} is assumed to be the relevant process, which was shown by kinetic simulations only recently \citep{2008ApJ...682L...5S}."
+ In the present work we locus on the main properues of the plasma instabihties and describe the influence of the fundamental mass-ratio z5/nm. in mixed electron-positron-proton compositions by means of relativistic three-dimensional simulations of counterstreaming plasmas., In the present work we focus on the main properties of the plasma instabilities and describe the influence of the fundamental mass-ratio $m_p/m_e$ in mixed electron-positron-proton compositions by means of relativistic three-dimensional simulations of counterstreaming plasmas.
+" The paper is organized as follows: In chapter 2. the underlying code is described briefly, in section 3 we illustrate the setup of the performed simulations."," The paper is organized as follows: In chapter \ref{description} the underlying code is described briefly, in section \ref{setup} we illustrate the setup of the performed simulations."
+ In chapter 4. we present the results of our simulations which we discuss and draw some conclusion in section 5.., In chapter \ref{results} we present the results of our simulations which we discuss and draw some conclusion in section \ref{discussion}.
+ Particle-1n-Cell simulations are an essential tool in understanding relativisuc collisionless plasma physics., Particle-in-Cell simulations are an essential tool in understanding relativistic collisionless plasma physics.
+" Therefore we developed a three-dimensional fully relativistic MPI-parallelised PiC code called ACRONYM (Annother Coode for moving Reelativisuc Obbjects, Noow with Yeee lattice and Maacroparticles)."," Therefore we developed a three-dimensional fully relativistic MPI-parallelised PiC code called ACRONYM nother ode for moving elativistic bjects, ow with ee lattice and acroparticles)."
+ Maxwells equations are evolved in ime by employing a second-order leapfrog scheme (see e.g. ?))., Maxwells equations are evolved in time by employing a second-order leapfrog scheme (see e.g. \citealt{Taflove}) ).
+ The particles aflect the electromagnetic fields through charge currents which are deposited on the grid by using a second-order Triangular Shaped Cloud (TSC) scheme (see e.g. ?)) adopted from ?.. , The particles affect the electromagnetic fields through charge currents which are deposited on the grid by using a second-order Triangular Shaped Cloud (TSC) scheme (see e.g. \citealt{1988csup.book.....H}) ) adopted from \citet{2001CoPhC.135..144E}. .
+The particles are moved via a force interpolation within the Boris push (2).., The particles are moved via a second-order force interpolation within the Boris push \citep{Boris1970}.
+" In order to guarantee that the divergence of the magnetic field remains close to zero, the electric and magnetic fields are stored in the form of a staggered grid, the so-called Yee-latlice (2).."," In order to guarantee that the divergence of the magnetic field remains close to zero, the electric and magnetic fields are stored in the form of a staggered grid, the so-called Yee-lattice \citep{1966ITAP...14..302Y}."
+ With this setup the code is second-order both in space and time., With this setup the code is second-order both in space and time.
+ Extensive tests of the code have been successfully completed from which we conclude that the total relative error in energy conservation is less that 3x107? and the divergence of the magnetic field stays below a value vB/B«107fp in the simulated space for all umes., Extensive tests of the code have been successfully completed from which we conclude that the total relative error in energy conservation is less that $3\ttt{-5}$ and the divergence of the magnetic field stays below a value $\left| \nabla \vec{B}/B \right| < 10^{-12}/\lambda_D$ in the simulated space for all times.
+" In the simulations presented here we use two counterstreaming plasma populations, one representing the background medium consisting of 6e and 6p* per cell (proton stream) and theother incorporating the"," In the simulations presented here we use two counterstreaming plasma populations, one representing the background medium consisting of $6
+e^-$ and $6 p^+$ per cell (proton stream) and theother incorporating the"
+moving up and down relative oa local continuum.,moving up and down relative to a local continuum.
+ We found similar correlations between he RY of the central emission and the residual intensities o ‘the τος ancl blue absorptions., We found similar correlations between the RV of the central emission and the residual intensities of the red and blue absorptions.
+ Apparently all the changes «observed. in the wings al velocities ος 100kms+ and ez kims Lowe due tothe variations of the centra reversal., Apparently all the changes observed in the wings at velocities $v \leq -$ $~{\rm km}~{\rm s}^{-1}$ and $v \geq$ $~{\rm km}~{\rm s}^{-1}$ are due to the variations of the central reversal.
+ We do not anticipate such large RY variations deep in the atmosphere where these wings are formed., We do not anticipate such large RV variations deep in the atmosphere where these wings are formed.
+ The conclusion is that the overall shape of the is determined by he central emission., The conclusion is that the overall shape of the is determined by the central emission.
+ We have performed. a period search. of the integrated equivalent width (IW) of he data set and. found a maximum in power at frequency 0.22 , We have performed a period search of the integrated equivalent width (EW) of the data set and found a maximum in power at frequency 0.22 $^{-1}$.
+Figures 1l and 12 show the phase diagram. [or he period 4.41 davs and a grev-scale representation of the phase spectrum. respectively.," Figures 11 and 12 show the phase diagram for the period 4.41 days and a grey-scale representation of the phase spectrum, respectively."
+ A phase spectrurn represents a two-dimensional case Of the algorithm where each velocity bin of the is treated as a time series of the intensity., A phase spectrum represents a two-dimensional case of the algorithm where each velocity bin of the is treated as a time series of the intensity.
+ Our target belongs to the group of stars for which the existence of a compact companion has heen proposed. in he literature., Our target belongs to the group of stars for which the existence of a compact companion has been proposed in the literature.
+ The task of disproving or confirming the xnary nature of the system can be tackled only if sullicientlv accurate analvsed observational data are available., The task of disproving or confirming the binary nature of the system can be tackled only if sufficiently accurate analysed observational data are available.
+ In this oper we used state-of-the-art. models. of atmospheres to determine the fundamental parameters of 111) 1882Q0., In this paper we used state-of-the-art models of atmospheres to determine the fundamental parameters of HD 188209.
+ To establish the presence of a possible companion we jwe studied the RVs of absorption lines by combining them, To establish the presence of a possible companion we have studied the RVs of absorption lines by combining them
+Due to tidal evolution. in the cistaut past the Moon's orbit was smaller.,"Due to tidal evolution, in the distant past the Moon's orbit was smaller."
+ As the orbital distatice is decreased. the Moou's orbital speed rises.," As the orbital distance is decreased, the Moon's orbital speed rises."
+" This increases the impact speeds on the leacliug hemisphere and makes ""catching up” to the Moon from behixl more difficult.", This increases the impact speeds on the leading hemisphere and makes “catching up” to the Moon from behind more difficult.
+ As Eq., As Eq.
+ 2. suggests. the degree of apex/antapex asymanetry on the Moon is expeced to increase as tlie orbital speec of the satellite does.," \ref{eq:alf} suggests, the degree of apex/antapex asymmetry on the Moon is expected to increase as the orbital speed of the satellite does."
+ We examined this by runuine other simulaious with au Earth-Moou separation of « =50. 38. 30. 20. and LOR.," We examined this by running other simulations with an Earth-Moon separation of $a = $ 50, 38, 30, 20, and 10 $R_\oplus$."
+ Figure 16. shows our results., Figure \ref{fig:dist} shows our results.
+ Assuming the same impactor orbital distribution in the past. ouly a wuld increase in the apex eihaucement is seen (since the orbital speed ouly increases. as Enνα).," Assuming the same impactor orbital distribution in the past, only a mild increase in the apex enhancement is seen (since the orbital speed only increases as $1/\sqrt(a)$ )."
+"- The increase. is. that expected based on the resulting. chauge iun a caused by the larger 655, (see Eq. 2)).", The increase is that expected based on the resulting change in $\alpha$ caused by the larger $v_{orb}$ (see Eq. \ref{eq:alf}) ).
+ As discussed in Sec. 2.3..," As discussed in Sec. \ref{sec-nearfar},"
+ the asymunetry between the near aud far hemisplieres should depeud ou the lunar distauce., the asymmetry between the near and far hemispheres should depend on the lunar distance.
+ Figure 16 shows the ratio between nearsice and farside craters from our simulations., Figure \ref{fig:dist} shows the ratio between nearside and farside craters from our simulations.
+ For all lunar distauces we lind very little asvinmetry., For all lunar distances we find very little asymmetry.
+ Thus. our results do uot support," Thus, our results do not support"
+simply for the velocity.,simply for the velocity.
+ We thus consider power spectra of both quantities., We thus consider power spectra of both quantities.
+ We have computed the power spectra for each code directly from gridded data using the same analysis script for both codes., We have computed the power spectra for each code directly from gridded data using the same analysis script for both codes.
+" For the SPH code, this means that the results have first been interpolated to grids of size 256°, 512? and 512? cells for the 128?, 256?, 512? particle calculations respectively."," For the SPH code, this means that the results have first been interpolated to grids of size $256^{3}$ , $512^{3}$ and $512^{3}$ cells for the $128^{3}$, $256^{3}$, $512^{3}$ particle calculations respectively."
+" The interpolation has been performed using a routine from the visualisation code, the details of which are described in ?.."," The interpolation has been performed using a routine from the visualisation code, the details of which are described in \citet{splashpaper}."
+ The main disadvantage of interpolating SPH data to a grid is that any resolution in the SPH code on scales smaller than the grid scale is lost., The main disadvantage of interpolating SPH data to a grid is that any resolution in the SPH code on scales smaller than the grid scale is lost.
+ This is most obvious when comparing quantities such as the maximum density on the particle data compared to the maximum density on the interpolated grid., This is most obvious when comparing quantities such as the maximum density on the particle data compared to the maximum density on the interpolated grid.
+" For example the maximum density interpolated onto a 512? grid (max(p/po)~3—4x 10?) is a factor of ~3—4 lower than the maximum density from the particles (max(p/po)~1—2x 10%) for snapshots from the 512? SPH calculation, which would remove some of the information from the high density tail of the PDF as discussed above."," For example the maximum density interpolated onto a $512^{3}$ grid $\max (\rho/\rho_{0}) \sim 3-4 \times 10^{2}$ ) is a factor of $\sim 3-4$ lower than the maximum density from the particles $\max (\rho/\rho_{0}) \sim 1-2 \times 10^{3}$ ) for snapshots from the $512^{3}$ SPH calculation, which would remove some of the information from the high density tail of the PDF as discussed above."
+ Whether or not the interpolation procedure affects the power spectrum calculation can be determined simply by comparing the results from interpolations to different sized grids., Whether or not the interpolation procedure affects the power spectrum calculation can be determined simply by comparing the results from interpolations to different sized grids.
+" We find that for the power spectrum, as one might expect given that it is a volumetric measurement, the power spectra are identical for different grid sizes apart from k’s very close to the grid scale."," We find that for the power spectrum, as one might expect given that it is a volumetric measurement, the power spectra are identical for different grid sizes apart from $k$ 's very close to the grid scale."
+" Power spectra for an arbitrary vector field w are constructed from the 3D Fourier transform, Theangle-averaged power spectrum was then obtained by the standard procedure of summing (asine.g.?).."," Power spectra for an arbitrary vector field ${\bf w}$ are constructed from the 3D Fourier transform, Theangle-averaged power spectrum was then obtained by the standard procedure of summing \citep[as in e.g.][]{kitsionasetal09}."
+ 9 , \ref{fig:ekin} 
+Nταν baud we have detected. simultancously aud with the sale instrunuents both the svuchrotron an the inverse Compton cluission in the spectrum of a blazar.,"X–ray band we have detected, simultaneously and with the same instruments, both the synchrotron and the inverse Compton emission in the spectrum of a blazar."
+ Simultaneous detection of both components in the X-ray spectra of blazar have already been reported by Ixubo et al. (, Simultaneous detection of both components in the X-ray spectrum of blazar have already been reported by Kubo et al. (
+1998) aud by Cüonual et al. (,1998) and by Giommi et al. (
+1999) for 85 0716|711. although not as clearly as in 2231.,"1999) for S5 0716+714, although not as clearly as in 231."
+ Iun the same feure we plot other three sets of quasismiultaueous observations: 1) the data of the 1996 multiwavelength campaign as reporte by Maisack et al. (, In the same figure we plot other three sets of quasi–simultaneous observations: i) the data of the 1996 multiwavelength campaign as reported by Maisack et al. (
+1997): i) the optical aud 5. ταν observations duriug 1995. when the source reached the brightest state iu the EGRET band: ii) the iufrared. optical. Nrav aud 5 raw data during 19911992. when the source was first detected iu the 5 rav baud bv ECRET.,"1997); ii) the optical and $\gamma$ –ray observations during 1995, when the source reached the brightest state in the EGRET band; iii) the infrared, optical, X–ray and $\gamma$ –ray data during 1991–1992, when the source was first detected in the $\gamma$ –ray band by EGRET."
+ The last two sets of data are not strictly simultaneous (the ~ rav fluxes detected in 19911992 refer to the sui of various »oiutines). but cau illustrate the different states of the source.," The last two sets of data are not strictly simultaneous (the $\gamma$ –ray fluxes detected in 1991–1992 refer to the sum of various pointings), but can illustrate the different states of the source."
+ The source was also detected by IRAS (Bhupey Neugebauer 1988). aud the corresponding IR fixes are reported in rofüe:sed.. even if they are not smüultaueous with any other observations.," The source was also detected by IRAS (Impey Neugebauer 1988), and the corresponding IR fluxes are reported in \\ref{fig:sed}, even if they are not simultaneous with any other observations."
+ Note that there are some mconsistenceies between the data in 19911992 as reported in Table 1 of vou Moutiguy et al. (, Note that there are some inconsistencies between the data in 1991–1992 as reported in Table 1 of von Montigny et al. (
+1995) aud the fluxes reported in Fig.,1995) and the fluxes reported in Fig.
+ 5 of the sale paper. which are consisteut with the flux reported o» Sreckumar et al. (," 5 of the same paper, which are consistent with the flux reported by Sreekumar et al. ("
+1996).,1996).
+ We lave reporte the data as shown in Fie., We have reported the data as shown in Fig.
+ 5 of von Moutieuy ct al. (, 5 of von Montigny et al. (
+1995).,1995).
+ As can ve seen. the19911992 7. rav spectrum is extremely hard (a~ (LAO. ," As can be seen, the1991–1992 $\gamma$ –ray spectrum is extremely hard $\alpha\sim 0.4\pm 0.4$ )."
+We coul uot fud the spectral iudex or the 1995 5 rav flux. but the shape of the spectra combining all observations together roni 1991 to 1995) Ίνα9A73+0.08 (ανα e al.," We could not find the spectral index for the 1995 $\gamma$ –ray flux, but the shape of the spectrum combining all observations together (from 1991 to 1995) is $\alpha\sim 0.73 \pm 0.18$ (Hartman et al."
+ 1999). sugsestiug hat the combined spectrum is steeper than he 19911992 spectrum (and suggesting that the 1995 spectrum is steeper still).," 1999), suggesting that the combined spectrum is steeper than the 1991--1992 spectrum (and suggesting that the 1995 spectrum is steeper still)."
+ Also shown are the upper ΕΕ iu the TeV uud. as derived by WIIIPPLE and ΠΕΕΑ observations (Maisack ot al.," Also shown are the upper limits in the TeV band, as derived by WHIPPLE and HEGRA observations (Maisack et al."
+ 1997) duxiug JanFeb 1996., 1997) during Jan–Feb 1996.
+ Alike other blazars. the SED of ON 231 is characterized w two broad compoucuts. the first peaking a IR.optical requenceies aud the secoud in the 5 rav baud.," Alike other blazars, the SED of ON 231 is characterized by two broad components, the first peaking at IR–optical frequencies and the second in the $\gamma$ –ray band."
+ The first is believed. to be svuchrotron cussion by a relativistic jet. while the secoud compoucut las been interpreted as svuchrotron self.Compton scattering. possibly iucludiug sole contribution from seed photons produced externally o the jet (see c.g. Ghlisellinuà ALadau 1996 aud reference herein) or svuchrotron by ultrarelativistic electronoositron pairs generated by relativistic protous (the xoton blazar model. Maunuheuu 1993. see \laisack οἳ al.," The first is believed to be synchrotron emission by a relativistic jet, while the second component has been interpreted as synchrotron self–Compton scattering, possibly including some contribution from seed photons produced externally to the jet (see e.g. Ghisellini Madau 1996 and reference therein) or synchrotron by ultra–relativistic electron--positron pairs generated by relativistic protons (the proton blazar model, Mannheim 1993, see Maisack et al."
+ 1997 for the application of this model to ON 231)., 1997 for the application of this model to ON 231).
+ The ouly other spectral information iu the X.ray band is from ROSAT: Comastri et al. (, The only other spectral information in the X–ray band is from ROSAT: Comastri et al. (
+1997). musing a sinele power law model. found an euergv spectral iudex a).= 1240.05. in agreement with Lamer et al. (,"1997), using a single power law model, found an energy spectral index $\alpha_x=1.2\pm 0.05$ , in agreement with Lamer et al. ("
+1996).,1996).
+ An earlier determination of the spectral shape using Eistcin data resulted in an unconstrained spectral iudex. (Worrall Wilkes 1990)., An earlier determination of the spectral shape using Einstein data resulted in an unconstrained spectral index (Worrall Wilkes 1990).
+ The shape of the Xrav spectra at the time of the ROSAT observation secs cdiffereut from the one determined by BeppoSAX., The shape of the X–ray spectrum at the time of the ROSAT observation seems different from the one determined by $Beppo$ SAX.
+ This could be due both to he narrower spectral coverage of ROSAT aud to the fact that the source was in a weaker state., This could be due both to the narrower spectral coverage of ROSAT and to the fact that the source was in a weaker state.
+" The 0.12.5 keV ROSAT spectrum. which is fatter than that measured by us in the same energy band (see index D, in Table 2). could be due to the contribution. iu this enerev band. of the fat component that BeppoSAX sees above 1 τον in May. auc above 2.5 keV in June."," The 0.1–2.5 keV ROSAT spectrum, which is flatter than that measured by us in the same energy band (see index $\Gamma_1$ in Table 2), could be due to the contribution, in this energy band, of the flat component that $Beppo$ SAX sees above 4 keV in May and above 2.5 keV in June."
+ If the break moves at lower energy when the source is weaker. as sugeested by the two BeppoSAX observations. then during the ROSAT observation. when the source was more than a factor of Lweaker. the break should be iuside the ROSAT band. or even at lower cuereies.," If the break moves at lower energy when the source is weaker, as suggested by the two $Beppo$ SAX observations, then during the ROSAT observation, when the source was more than a factor of 4 weaker, the break should be inside the ROSAT band, or even at lower energies."
+ Tn this case the spectra detected by ROSAT would be cither a combination of svuchrotron and Compton oenüssion or purely due to Compton scattering. explaining the relative flatuess of the ROSAT spectrum.," In this case the spectrum detected by ROSAT would be either a combination of synchrotron and Compton emission or purely due to Compton scattering, explaining the relative flatness of the ROSAT spectrum."
+ A stecpening of the spectrum eoiue to softer energies in the SED is also required by the quasi siultauecous IR.optical data (Massaro et al., A steepening of the spectrum going to softer energies in the SED is also required by the quasi simultaneous IR–optical data (Massaro et al.
+ 1991)., 1994).
+ The observed “flare” iu the soft Norav baud is ποΊσα] (equal rise and decay timescale). sugeestiug hat the variabilitv timescale is determined bv the light Crosse time of the ciitting region. ο (see Chiaberge (κο 1999).," The observed “flare” in the soft X–ray band is symmetrical (equal rise and decay timescale), suggesting that the variability timescale is determined by the light crossing time of the emitting region, $R/c$ (see Chiaberge Ghisellini 1999)."
+ This in turi implies that the cooling time is shorter than Rc. allowing us to put limits on he value of the magnetic field and on the euergv of the electrous producing the variable flux at the oherved frequency 7.," This in turn implies that the cooling time is shorter than $R/c$, allowing us to put limits on the value of the magnetic field and on the energy of the electrons producing the variable flux at the oberved frequency $\nu_x$."
+" Using fear=5 hours aud p,=3<1019 DD. we derive DB(0.48’ Gauss aud 5,«L510051/7,"," Using $t_{\rm var}=5$ hours and $\nu_x=3\times 10^{16}$ Hz, we derive $B>0.4\delta^{-1/3}$ Gauss and $\gamma_x< 1.5\times 10^5\delta^{-1/3}$."
+ Tere 8=[TVI?1eos0|3 isthe Doppler beaming factor. where 0 is the viewing angle.," Here $\delta=[\Gamma-\sqrt{\Gamma^2-1}\cos\theta]^{-1}$ is the Doppler beaming factor, where $\theta$ is the viewing angle."
+" Since the peak of the svuchrotron enissjon nmst be at a frequency vw<3o.10! IIz. the correspouding Loreutz factor of the clectrous cinitting at the peak is 5,<1.5«10!à.13)"," Since the peak of the synchrotron emission must be at a frequency $\nu_{\rm s} < 3\times 10^{14}$ Hz, the corresponding Lorentz factor of the electrons emitting at the peak is $\gamma_{\rm s}< 1.5\times 10^4\delta^{-1/3}$."
+" This implies that the peak of the self Compton fux is at iinτο)Ίος«3708. ATV, Tighter constraints can be obtained asstning that the optical and Nray cinissiou are cospatial. and that also the cooling time of the optical cuuitting clectrous is shorter than the light crossing time."," This implies that the peak of the self Compton flux is at $h\nu_{\rm c} = (4/3)\gamma^2_{\rm s}h\nu_{\rm s} < 370 \delta^{-2/3}$ MeV. Tighter constraints can be obtained assuming that the optical and X–ray emission are cospatial, and that also the cooling time of the optical emitting electrons is shorter than the light crossing time."
+ Massaro ct al. (, Massaro et al. (
+1999) report intranight variations iu the T aud D bauds of 0.2 magnitude in 1l hour. with an approximately sviuunetric time profile.,"1999) report intranight variations in the I and B bands of 0.2 magnitude in 1 hour, with an approximately symmetric time profile."
+" ο again fü,=5 hours we obtain D»1.5813 Cass and x.9.7«1078. 1/75, "," Assuming again $t_{\rm var}=5$ hours, we obtain $B>1.5\delta^{-1/3}$ Gauss and $\gamma_{\rm s}< 9.7\times 10^3\delta^{-1/3}$ ."
+The far IR τος ray eissiou from blazars cau be explained by simple homogeneous aud onezone svuchrotron inverse Compton models. with the enüttius region movie relativistically towards the observer.," The far IR to $\gamma$ –ray emission from blazars can be explained by simple homogeneous and one–zone synchrotron inverse Compton models, with the emitting region moving relativistically towards the observer."
+ The inverse Comptou, The inverse Compton
+" 0.74, 1.1 and 1.5Jy at 1.4 mm for the flux of the nucleus and the two radio lobes which is in agreement with what we obtain.","$0.74$ , $1.1$ and $1.5\,\mathrm{Jy}$ at 1.4 mm for the flux of the nucleus and the two radio lobes which is in agreement with what we obtain."
+" shows a secure detection of NGC1068, a nearby galaxy with an active galactic nucleus."," shows a secure detection of NGC1068, a nearby galaxy with an active galactic nucleus."
+ The map is obtained with 5 scans and a total integration time of 1260 seconds., The map is obtained with 5 scans and a total integration time of 1260 seconds.
+ Its flux is 142+25mJy at 1.4 mm and 66+3mJy at 2 mm.," Its flux is $142\pm25\,\mathrm{mJy}$ at 1.4 mm and $66\pm3\,\mathrm{mJy}$ at 2 mm."
+" This is the central flux measured with a Gaussian of 12 and 19 arcseconds, respectively."," This is the central flux measured with a Gaussian of 12 and 19 arcseconds, respectively."
+" For this map, a sky noise decorrelation has been used, which is based on a linear regression with the detector signals when off-source."," For this map, a sky noise decorrelation has been used, which is based on a linear regression with the detector signals when off-source."
+ The North East - South West extension at 1mm is larger but aligned with interferometric IRAM Plateau de Bure Interferometer (PdBI) observations (Kripsetal. 2006))., The North East - South West extension at 1mm is larger but aligned with interferometric IRAM Plateau de Bure Interferometer (PdBI) observations \cite{Krips2006}) ).
+ The PdBI measured flux for the core and the jet is 28+2mJy at 231GHz.," The PdBI measured flux for the core and the jet is $28\pm2\,\mathrm{mJy}$ at $231\,\mathrm{GHz}$."
+" This is smaller than the NIKA flux measurement which is more consistent with the flux of 170+30mJy measured by Thronsonetal.1987,, indicating the presence of a diffuse extended component."," This is smaller than the NIKA flux measurement which is more consistent with the flux of $170\pm30\,\mathrm{mJy}$ measured by \cite{Thronson1987}, indicating the presence of a diffuse extended component."
+" This component is expected (Hildebrandetal. 1977)), based on consideration of the far infrared spectral energy distribution."," This component is expected \cite{Hildebrand1977}) ), based on consideration of the far infrared spectral energy distribution."
+ It is likely to come from heated dust in the circumnuclear region., It is likely to come from heated dust in the circumnuclear region.
+" From deep integration on weak sources and using sky noise decorrelation, we are able to derive the effective sensitivity of the camera, in the present early state of data reduction and sky noise subtraction techniques."," From deep integration on weak sources and using sky noise decorrelation, we are able to derive the effective sensitivity of the camera, in the present early state of data reduction and sky noise subtraction techniques."
+" We obtain a weak-source flux detectivity of 450 and 37,mJy-s!? at 1.4 and 2 mm respectively."," We obtain a weak-source flux detectivity of 450 and $37,\mathrm{mJy}\cdot
+s^{1/2}$ at 1.4 and 2 mm respectively."
+ The 1.4 mm detectivity is satisfactory for an initial 1.4 mm KID prototype., The 1.4 mm detectivity is satisfactory for an initial 1.4 mm KID prototype.
+ The 2 mm detectivity shows a major improvement by a factor 3 with the best value obtained in the 2009 NIKA run (Monfardinietal. 2010))., The 2 mm detectivity shows a major improvement by a factor 3 with the best value obtained in the 2009 NIKA run \cite{Monfardini:29}) ).
+ This detectivity is almost at the level of the state-of-art APEX-SZ TES detectors (Schwanetal. 2010)) albeit obtained here with a larger telescope., This detectivity is almost at the level of the state-of-art APEX-SZ TES detectors \cite{Schwan2010}) ) albeit obtained here with a larger telescope.
+ An NET of 6mK-s!? is deduced from this value., An NET of $6\mathrm{ mK} \cdot s^{1/2}$ is deduced from this value.
+" A sensitivity to the SZ effect can be obtained as 3-10?Ar!? in the y Compton parameter for onebeam, although this may be complicated by the SZ extension."," A sensitivity to the SZ effect can be obtained as $3 \cdot
+10^{-5} hr^{1/2}$ in the $y$ Compton parameter for onebeam, although this may be complicated by the SZ extension."
+ SZ measurements will be reported in a later publication., SZ measurements will be reported in a later publication.
+(VE-rfNSM)Te=MossTVVMSτονPON.Ne.EToNa). where a.b.e are integers.,"$\langle N_1^a N_2^b N_3^c \rangle = 
+\sum_{N_1,N_2,N_3}
+N_1^aN_2^bN_3^c P(N_1,N_2,N_3)$, where $a,b,c$ are integers."
+. In particular.. (Vj)- is (he average population size of the specie A; on a grain.," In particular, $\langle N_{i}\rangle$ is the average population size of the specie $X_i$ on a grain."
+" The production rate per grain. R(X) (molecules ο). of Xj molecules produced by the reaction JV;4-X;—Ny is given by ΜΑ=CA; A;,)CV;IN). or by REX.)=AGCV;CV;—1)) in case that 7—j."," The production rate per grain, $R(X_k)$ (molecules $^{-1}$ ), of $X_k$ molecules produced by the reaction $X_i + X_j \rightarrow X_k$ is given by $R(X_k) = (A_i+A_j) \langle N_i \ N_j \rangle$, or by $R(X_k)=A_i \langle N_i(N_i-1) \rangle$ in case that $i=j$."
+ In numerical simulations the master equation must be (truncated in order to keep the nunbber of equations finite., In numerical simulations the master equation must be truncated in order to keep the number of equations finite.
+ This can be done by setting upper cutolfs VMS. 7=Leo.../ on the population sizes. where J is (he number of reactive species.," This can be done by setting upper cutoffs $N_i^{\max}$, $i=1,\dots,J$ on the population sizes, where $J$ is the number of reactive species."
+" However. the number of coupled equations. Ny=IL,CN+1). grows exponentially with (he number of reactive species."," However, the number of coupled equations, $N_E = \prod_{i=1}^J (N_i^{\max}+1)$, grows exponentially with the number of reactive species."
+ This severely limits the applicability of the master equation to interstellar chemistry (Slantcheva&Herbst2003)., This severely limits the applicability of the master equation to interstellar chemistry \citep{Stantcheva2003}.
+. To reduce the number of equations one tries to use the lowest possible eutoffs under (he given conditions., To reduce the number of equations one tries to use the lowest possible cutoffs under the given conditions.
+ In any. case. to enable all reaction processes {ο take place. the cutolls must satisly 2777>2 for species Chat form homonuclear diatonic molecules (110 Ου. ete.)," In any case, to enable all reaction processes to take place, the cutoffs must satisfy $N_i^{\rm max} \ge 2$ for species that form homonuclear diatomic molecules $_2$ $_2$, etc.)"
+ and NPS>] for other species., and $N_i^{\rm max} \ge 1$ for other species.
+ The average population sizes of the reactive species and the reaction rates are completely determined by all the first moments and selected second moments of the distribution PN.No.Aa).," The average population sizes of the reactive species and the reaction rates are completely determined by all the first moments and selected second moments of the distribution $P(N_1,N_2,N_3)$."
+ Therefore. a closed set of equations for the time derivatives of these first and second moments could provide complete information on the population sizes aud reaction rates.," Therefore, a closed set of equations for the time derivatives of these first and second moments could provide complete information on the population sizes and reaction rates."
+ For (he simple network considered here one needs equations for the time derivatives of the first moments (V1). (No) and CV4) and of the second moments (V7). (V3). ΝΑ) and. ΝΑ) [nodes and edges. respectively. in the graph shown in Fig. Lj].," For the simple network considered here one needs equations for the time derivatives of the first moments $\langle N_1 \rangle$, $\langle N_2 \rangle$ and $\langle N_3 \rangle$ and of the second moments $\langle N_1^2 \rangle$, $\langle N_2^2 \rangle$, $\langle N_1 N_2 \rangle$ and $\langle N_1 N_3 \rangle$ [nodes and edges, respectively, in the graph shown in Fig. \ref{fig:1}] ]."
+ Such equations are obtained by Lalàng the time derivative of each moment and using Eq. (1)), Such equations are obtained by taking the time derivative of each moment and using Eq. \ref{eq:Master}) )
+ to express (he lime derivatives of the probabilities (Lipshtat&Biba2003)., to express the time derivatives of the probabilities \citep{Lipshtat2003}.
+. ere we show two of the resulting moment equations: In these equations. the time derivative of each moment is expressed as a linear combination of several olher moments.," Here we show two of the resulting moment equations: In these equations, the time derivative of each moment is expressed as a linear combination of several other moments."
+ However. (he right hand sides of these equations include third order moments for which we have no equations.," However, the right hand sides of these equations include third order moments for which we have no equations."
+ In order to close the set of moment equations we must express the third order moments in terms of first and second order moments (Lipshtat&Diham2003)., In order to close the set of moment equations we must express the third order moments in terms of first and second order moments \citep{Lipshtat2003}.
+. This can be done bv imposing the following constraint on the master equation: al any given (nme. al most (wo atoms or molecules can be adsorbed simultaneously on the surface.," This can be done by imposing the following constraint on the master equation: at any given time, at most two atoms or molecules can be adsorbed simultaneously on the surface."
+ Furthermore. these (wo atoms or molecules must be from," Furthermore, these two atoms or molecules must be from"
+the surface brightness enhancement in the arms.,the surface brightness enhancement in the arms.
+ However. the pressure difference seems unphysical. since the time for the pressure to come to equilibrium (the sound crossing time of the arm) is short compared to the sound travel time along the length of the narrow arm.," However, the pressure difference seems unphysical, since the time for the pressure to come to equilibrium (the sound crossing time of the arm) is short compared to the sound travel time along the length of the narrow arm."
+ Therefore. the axis of the arm would need to make a sufficiently small angle with our line-of-sight (roughly 157)) to make its projected path length about 4 times its width. in order to account for its high emission measure.," Therefore, the axis of the arm would need to make a sufficiently small angle with our line-of-sight (roughly ) to make its projected path length about 4 times its width, in order to account for its high emission measure."
+ This seems unlikely. given that the brightening along the arm is at a comparable radius to that of the eastern arm and both are at a radius similar to that of the 14 kpe (3) ring.," This seems unlikely, given that the brightening along the arm is at a comparable radius to that of the eastern arm and both are at a radius similar to that of the 14 kpc $3'$ ) ring."
+ However. in computing the overpressure of the arm. we did not include any contribution to the external pressure from non-thermal particles or from magnetic field which could contribute to the confinement.," However, in computing the overpressure of the arm, we did not include any contribution to the external pressure from non-thermal particles or from magnetic field which could contribute to the confinement."
+ Magnetic tension could also serve to confine the gas in the southwestern arm since the radio emission from the southwestern arm appears to spiral around the X-ray gas (see Fig. 11))., Magnetic tension could also serve to confine the gas in the southwestern arm since the radio emission from the southwestern arm appears to spiral around the X-ray gas (see Fig. \ref{fig:overlay}) ).
+ Alternatively. the southwestern arm may not be formed by rising bubbles as seen in the eastern arm. but is instead a thin sheath of gas exterior to a large plasma bubble related to the southwestern radio arm.," Alternatively, the southwestern arm may not be formed by rising bubbles as seen in the eastern arm, but is instead a thin sheath of gas exterior to a large plasma bubble related to the southwestern radio arm."
+ This is consistent with a roughly increase in surface brightness from west to east across the arm., This is consistent with a roughly increase in surface brightness from west to east across the arm.
+ We have presented a discussion of several of the remarkable structures seen in the Chandra. XMM-Newton. and ROSAT HRI observations of M87.," We have presented a discussion of several of the remarkable structures seen in the Chandra, XMM-Newton, and ROSAT HRI observations of M87."
+ Many of these. particularly the bubbles emanating from the central region. the nearly circular rings of enhanced emission at 14 kpe and 17 kpe. and the brightening of the X-ray arms at these radit. can be attributed to AGN outbursts.," Many of these, particularly the bubbles emanating from the central region, the nearly circular rings of enhanced emission at 14 kpc and 17 kpc, and the brightening of the X-ray arms at these radii, can be attributed to AGN outbursts."
+ The 14 kpe and 17 kpe rings. similar to the “ripples” seen in the Perseus cluster (Fabian et al.," The 14 kpc and 17 kpc rings, similar to the “ripples” seen in the Perseus cluster (Fabian et al."
+ 2003). can be interpreted as shock waves driven by the current outburst that began about 10 years ago.," 2003), can be interpreted as shock waves driven by the current outburst that began about $10^7$ years ago."
+ The outburst also inflated the inner radio lobes (and cocoon)., The outburst also inflated the inner radio lobes (and cocoon).
+Outbursts like those that produced these shocks can quench the M87 cooling flow. if they occur approximately every 3«107. years.," Outbursts like those that produced these shocks can quench the M87 cooling flow, if they occur approximately every $3\times10^7$ years."
+" Since the enthalpy associated with the inner cavities. produced by the outburst. is only of the energy of the outburst. shock heating is probably the dominant heating mechanism for the gas in the inner regions of ""cooling flow"" systems."," Since the enthalpy associated with the inner cavities, produced by the outburst, is only of the energy of the outburst, shock heating is probably the dominant heating mechanism for the gas in the inner regions of “cooling flow” systems."
+ At larger radii. we see highly enriched gas along the outer edge of the southern radio lobe (the 37 kpe arc).," At larger radii, we see highly enriched gas along the outer edge of the southern radio lobe (the 37 kpc arc)."
+ Asymmetric gas distributions at radit of ~50 kpe may be evidence for older C»107 years) outbursts as are the outer radio lobes., Asymmetric gas distributions at radii of $\sim$ 50 kpc may be evidence for older $>10^8$ years) outbursts as are the outer radio lobes.
+ The hot X-ray emitting gas contains reflections of previous episodes of AGN activity in the form of bubbles and their bright rims. shocks. and buoyantly uplifted gas structures.," The hot X-ray emitting gas contains reflections of previous episodes of AGN activity in the form of bubbles and their bright rims, shocks, and buoyantly uplifted gas structures."
+ With the detailed observations at X-ray and radio wavelengths of M87. we can probe the interaction between the central AGN. the relativistic plasma. and the X-ray gas.," With the detailed observations at X-ray and radio wavelengths of M87, we can probe the interaction between the central AGN, the relativistic plasma, and the X-ray gas."
+ We are beginning to understand the cyclic heating of the X-ray gas and the energy transfer mechanisms between the central supermassive black hole and the hot gaseous atmosphere that surrounds central cluster galaxies., We are beginning to understand the cyclic heating of the X-ray gas and the energy transfer mechanisms between the central supermassive black hole and the hot gaseous atmosphere that surrounds central cluster galaxies.
+ We acknowledge stimulating discussions with D. Harris. L. David. M. Begelman. and R.. Sunyaev.," We acknowledge stimulating discussions with D. Harris, L. David, M. Begelman, and R., Sunyaev."
+ This work was supported by NASA contracts NAS8-38248. NAS8-01130. NAS8-030060. the Chandra Science Center. the Smithsonian Institution. and MPI fürr Astrophysik.," This work was supported by NASA contracts NAS8-38248, NAS8-01130, NAS8-03060, the Chandra Science Center, the Smithsonian Institution, and MPI fürr Astrophysik."
+starbursts itis in the disks of normal galaxies. where help would be needed.,"starbursts – it is in the disks of normal galaxies, where help would be needed."
+ If the mid-Ilt is to be useful as à star-formation rate (SER) indicator. calibrators with the other methods. are necessary due to the complexity of theoretically deriving an accurate relation between mid-LHt emission ancl amount of voung stars.," If the mid-IR is to be useful as a star-formation rate (SFR) indicator, calibrators with the other methods are necessary due to the complexity of theoretically deriving an accurate relation between mid-IR emission and amount of young stars."
+ Indeed. Ho. emission has been shown to correlate with. micl-LR luminosity of spiral ealaxies (Roussel ct al.," Indeed, $\alpha$ emission has been shown to correlate with mid-IR luminosity of spiral galaxies (Roussel et al."
+ 2001b. Vigroux et al.," 2001b, Vigroux et al."
+ 1999. Cesarsky Sauvage 1999).," 1999, Cesarsky Sauvage 1999)."
+ Also [ar-L1t seems to correlate linearly with mic-Hand Ho if only disks are considered. and SEI. can thus be estimated (Vigroux et al 1999. Roussel 20015).," Also far-IR seems to correlate linearly with mid-IRand $\alpha$ if only disks are considered, and SFR can thus be estimated (Vigroux et al 1999, Roussel 2001b)."
+ The relations do not hold in regions of more intense star-formation nnucleus). and thus nuclear star-formation could. confuse a global SER determination.," The relations do not hold in regions of more intense star-formation nucleus), and thus nuclear star-formation could confuse a global SFR determination."
+ Vigroux ct aancl Roussel argue that this precisely is the reason for non-linearity in global far-LRo Lda relations., Vigroux et and Roussel argue that this precisely is the reason for non-linearity in global far-IR $\alpha$ relations.
+ Thus. only limits toSER can be calculated from integrated mid-Lt luminosity. and the need for information on the proportions of disk and nuclear Hi-emission is high-lighted.," Thus, only limits to SFR can be calculated from integrated mid-IR luminosity, and the need for information on the proportions of disk and nuclear IR-emission is high-lighted."
+ The micl-L Hux ratio is helpful in tracing the star-forming activity. as discussed. belore.," The mid-IR flux ratio is helpful in tracing the star-forming activity, as discussed before."
+ We also expect. the near-IR to add to the information., We also expect the near-IR to add to the information.
+ To illustrate these effects. lig.," To illustrate these effects, Fig."
+ 12 shows all those galaxies from the comparison sample discussed in Sections 3.3. and 4.l.. which had IAS fluxes available.," \ref{fircols} shows all those galaxies from the comparison sample discussed in Sections \ref{compsamp} and \ref{class}, which had fluxes available."
+ The upper left panel shows an, The upper left panel shows an
+"Behi&Freeman2003)). v Dellazzuvietal.2008:Carrettaet201023). Carrettaetal.(901010) all strong interaction with the Galactic main body triggeredao a chain of events in the so called preeersor of curveut CCS, whose final products are the svstenis we are seciusc» (see Carrettaetal.9010 for a detailed description).","\citealt{bek03}) $\omega$ \citealt{bok09,bel08,car10a}) \cite{car10b} $all$ strong interaction with the Galactic main body triggered a chain of events in the so called $precursor$ of current GCs, whose final products are the systems we are seeing (see \citealt{car10b} for a detailed description)."
+ The same imechanisz may also work within a singlec» dSph. as in the case of Fornax or Sagittarius.," The same mechanism may also work within a single dSph, as in the case of Fornax or Sagittarius."
+ This scenario. although not eutirelv new (e.g... S¢sarleZinn 1978)) may explain several characteristic*oof GCs. includiug the multiple stellar generations. found iu all objects iuvestigated so far (Carrettaetal.200a.b and Cattonetal.200L for a review).," This scenario, although not entirely new (e.g., \citealt{sea78}) ), may explain several characteristics of GCs, including the multiple stellar generations, found in $all$ objects investigated so far \citealt{car09a,car09b} and \citealt{gra04} for a review)."
+ If is currvutly well assessed that the seconud-generation stars (prexvutly constituting the bulk of the cluster population) should have formed from the ejecta of only a fraction of the first-eeneration (primordial) stars (c.g. Cawattonctal.2001:Prautzos&Charbonuel 2006)).," It is currently well assessed that the second-generation stars (presently constituting the bulk of the cluster population) should have formed from the ejecta of only a fraction of the first-generation (primordial) stars (e.g., \citealt{gra01,pra06}) )."
+ To account for the present elieniical composition. a precursor barvondc 1Lass Hyout 20-50 times larecr than the current mass of the GCs is required (Carrettaal.20105).," To account for the present chemical composition, a precursor baryonic mass about 20-50 times larger than the current mass of the GCs is required \citep{car10b}."
+. However. the proposed scenario is still quitative: for instance. it is not obvious that only one GC should be the final out]it of i individual precursor.," However, the proposed scenario is still qualitative; for instance, it is not obvious that only $one$ GC should be the final output of an individual precursor."
+ Examples ofbinary aud iiultiple (‘listers are frequently observed in the Large Magellauic Cloud (e... Dieballetal. 2002)).," Examples of binary and multiple clusters are frequently observed in the Large Magellanic Cloud (e.g., \citealt{die02}) )."
+ To clarify these issues. we added NCC 1851 to our ongoing FLAMES survey studving the NaO anticorrelation in GCs (Carrettaetal.2006).," To clarify these issues, we added NGC 1851 to our ongoing FLAMES survey studying the Na-O anticorrelation in GCs \citep{car06}."
+ NCC |851 has a bimodal horizoutal brauch (IB) and several other peculiarities., NGC 1851 has a bimodal horizontal branch (HB) and several other peculiarities.
+ The color-magnitude diagram (CMD) shows a double subeiaut branch (SCB: Miloueοal. 2008))., The color-magnitude diagram (CMD) shows a double subgiant branch (SGB; \citealt{mil08}) ).
+ Two distinct red eiant brauch. (ROB) sequelicoCwRs were discovered by Leeetal.(2009a) with Cauweby photometry. and confirmed by Tauoetal.(2019) using broad-band filters.," Two distinct red giant branch (RGB) sequences were discovered by \cite{lee09a} with $uvby$ photometry, and confirmed by \cite{han09} using broad-band filters."
+" In Carettaotal.(320100) Wwe challenged the hypothesis advanced by Leeetal.(2091a) of Ca variatious in GCs, ascribed to a possible polhtion"," In \cite{car10c} we challenged the hypothesis advanced by \cite{lee09b} of Ca variations in GCs, ascribed to a possible pollution"
+and (he optical redshift values Crom Table 1) for sources that exhibit absorption only.,and the optical redshift values (from Table 1) for sources that exhibit absorption only.
+ Colhunns (3) ancl (4) give the logarithms of the FIR. and the total I. huninosities calculated using equations | and 2. respectively.," Columns (3) and (4) give the logarithms of the FIR and the total IR luminosities calculated using equations 1 and 2, respectively."
+ Column (5) is the total on-source integration Unies., Column (5) is the total on-source integration times.
+ Columns (6) lists the rms noise levels., Columns (6) lists the rms noise levels.
+ Column (7) lists the heliocentric velocities al the peak of the absorption feature in each source., Column (7) lists the heliocentric velocities at the peak of the absorption feature in each source.
+ Column (3) is the EWIIM velocity values of the overall observed absorplion profile., Column (8) is the FWHM velocity values of the overall observed absorption profile.
+ Column (9) lists the peak optical depths., Column (9) lists the peak optical depths.
+ Column (10) eives the column densities divided by the spin temperature which are derived assumius a covering factor of unity ancl using (he equation (e.g. Rohs1986)): We detected. OL] 18 cm main lines in 7 galaxies: 3 in emission (OIIMs) and 4 in absorption., Column (10) gives the column densities divided by the spin temperature which are derived assuming a covering factor of unity and using the equation (e.g. \citealp{Roh86}) ): We detected OH 18 cm main lines in 7 galaxies; 3 in emission (OHMs) and 4 in absorption.
+ Of these. 1 OIIM and 3 OL absorbers are new detectüions (Figure 3).," Of these, 1 OHM and 3 OH absorbers are new detections (Figure 3)."
+ For these spectra. (he 1667.359 MIIz line was used [or deriving the velociiv scale.," For these spectra, the 1667.359 MHz line was used for deriving the velocity scale."
+ There were no detections of the OIL 18 cm satellite lines al 1612 and 1720 MIIz in anv of the galaxies of our sample., There were no detections of the OH 18 cm satellite lines at 1612 and 1720 MHz in any of the galaxies of our sample.
+ Figure 4 shows the spectra of (he 3 sources wilh previously known OII 18 cm main lines., Figure 4 shows the spectra of the 3 sources with previously known OH 18 cm main lines.
+ As for the 21 em spectra in Figures 1 2. Figures 3 4 present the final spectra. bul without the subtraction of fitted polvnonmials.," As for the 21 cm spectra in Figures 1 2, Figures 3 4 present the final spectra, but without the subtraction of fitted polynomials."
+ Table 4 summarizes the parameters derived fom the OIM detections., Table 4 summarizes the parameters derived from the OHM detections.
+ The values in this table are derived using the baseline-subtracted spectra., The values in this table are derived using the baseline-subtracted spectra.
+ Coblunn (1) lists the IRAS names of the galaxies., Column (1) lists the IRAS names of the galaxies.
+ Column (2) is the (total on-source integration times., Column (2) is the total on-source integration times.
+ Column (3) gives the rms noise levels. and column (4) lists the heliocentric velocities at the center of the 1667.359 ΣΠ line emission.," Column (3) gives the rms noise levels, and column (4) lists the heliocentric velocities at the center of the 1667.359 MHz line emission."
+ Columns (5) and (3) are the peak flux densities of the 1667 and 1665 Ην lines respectively. columns (6) ancl (9) are their [ull velocity widths at half maximum. ancl columus (7) and (10) list their integrated. flux densities.," Columns (5) and (8) are the peak flux densities of the 1667 and 1665 MHz lines respectively, columns (6) and (9) are their full velocity widths at half maximum, and columns (7) and (10) list their integrated flux densities."
+ Column (11) lists the hyperfine ratios obtained by dividing the integrated flux density of the 1667 MIIz line bv that at 1665 Mllz for each source., Column (11) lists the hyperfine ratios obtained by dividing the integrated flux density of the 1667 MHz line by that at 1665 MHz for each source.
+ In thermodvnamic equilibrium conditions. (his ratio would be Ry=1.8. and il increases as the saturation of a masing region increases.," In thermodynamic equilibrium conditions, this ratio would be $R_H=1.8$, and it increases as the saturation of a masing region increases."
+ Column (12) lists the logarithms of the predicted OIL luminosities (in units of £..) which have been calculated using the following equation from Ixandalian(1996): while column (13) gives the logarithuis of the measured isotropic OII line luminosities. which represents (he combined integrated flux densities of (he two OII 18 cm main Ines.," Column (12) lists the logarithms of the predicted OH luminosities (in units of ${L_{\odot}}$ ) which have been calculated using the following equation from \citet{KAN96}: while column (13) gives the logarithms of the measured isotropic OH line luminosities, which represents the combined integrated flux densities of the two OH 18 cm main lines."
+ , 
+ t, 
+ tr, 
+ tra, 
+ tran, 
+ trans, 
+ transi, 
+ transit, 
+ transiti, 
+ transitio, 
+ transition, 
+ transitions, 
+ transitions , 
+ transitions N, 
+ transitions N., 
+ transitions N.O, 
+ transitions N.O., 
+ transitions N.O. , 
+ , 
+same class of SGRs and AXPs. then one would like to interpret the observations within the context of the magnetar model.,"same class of SGRs and AXPs, then one would like to interpret the observations within the context of the magnetar model."
+ The magnetar model as currently formulated does not make specific predictions for on/off states of the pulsars as well as for their IR fluxes., The magnetar model as currently formulated does not make specific predictions for on/off states of the pulsars as well as for their IR fluxes.
+ Variations in the quiescent emission have been observed following an X-ray burst in (Woods et al., Variations in the quiescent emission have been observed following an X–ray burst in (Woods et al.
+ 2003)., 2003).
+ Soon after the burst. the persistent X—ray flux washigher. the temperature also higher. and the blackbody radius much smaller than in the low quiescent state.," Soon after the burst, the persistent X--ray flux was, the temperature also higher, and the blackbody radius much smaller than in the low quiescent state."
+ Some of the pproperties in the quiescent emission before and after the burst are consistent (at least qualitatively). with the those of un its low and high state.," Some of the properties in the quiescent emission before and after the burst are consistent (at least qualitatively), with the those of in its low and high state."
+" Moreover. the slow decay of the ""high-state"" X-ray emission of iis found to be in the range of those of SGR 1627-41 and SGR 1900414 (Ibrahim et al."," Moreover, the slow decay of the “high-state” X–ray emission of is found to be in the range of those of SGR 1627-41 and SGR 1900+14 (Ibrahim et al."
+ 2003: If wwere an isolated object. then (a hypothetical) accretion would have to proceed through a fall-back disk (Chatterjee et al.," 2003; If were an isolated object, then (a hypothetical) accretion would have to proceed through a fall-back disk (Chatterjee et al."
+ 2000. Alpar 2001)," 2000, Alpar 2001)."
+ While short-term. small scale fluctuations are expected in this case. by analogy with most accreting neutron star systems. large scale variations might require special conditions.," While short-term, small scale fluctuations are expected in this case, by analogy with most accreting neutron star systems, large scale variations might require special conditions."
+ In particular. as the object spins down. it might occasionally switch from a propeller phase to an aceretion phase.," In particular, as the object spins down, it might occasionally switch from a propeller phase to an accretion phase."
+ During the propeller phase. accretion is inhibited. and the star should be bright in X-rays due to its thermal emission.," During the propeller phase, accretion is inhibited, and the star should be bright in X–rays due to its thermal emission."
+ Indeed. in its low state. hhad an X-ray luminosity of the order of a few «1045 erg/s. which is typical of other known thermal emitting sources (e.g. Becker Trümmper 1997).," Indeed, in its low state, had an X–ray luminosity of the order of a few $\times 10^{33} d_5$ erg/s, which is typical of other known thermal emitting sources (e.g. Becker Trümmper 1997)."
+ Moreover. in its low state the pulsed fraction of (af any) decreased considerably. suggesting emission from the whole surface of the star as expected in a cooling object (G03).," Moreover, in its low state the pulsed fraction of (if any) decreased considerably, suggesting emission from the whole surface of the star as expected in a cooling object (G03)."
+ Note that in a more general scenario. the CCOs would be AXPs in the above described state.," Note that in a more general scenario, the CCOs would be AXPs in the above described state."
+ If the high state is due to resumed accretion onto the magnetic poles. then larger pulsed fractions would be naturally explained.," If the high state is due to resumed accretion onto the magnetic poles, then larger pulsed fractions would be naturally explained."
+ The aceretion luminosity could be much higher than the thermal emission. and would dominate the emitted. spectrum.," The accretion luminosity could be much higher than the thermal emission, and would dominate the emitted spectrum."
+" These properties appear consistent with those ofXTEJ1810—197.. and the properties of the IR counterpart could be explained in. terms of a fall-back disk of size < a few «10!"" em (using the spectral models of Perna et al."," These properties appear consistent with those of, and the properties of the IR counterpart could be explained in terms of a fall-back disk of size $\la$ a few $\times 10^{10}$ cm (using the spectral models of Perna et al."
+ 2000 and Perna Hernquist 2000)., 2000 and Perna Hernquist 2000).
+ However. the X-ray flux decay reported by Ibrahim et al. (," However, the X–ray flux decay reported by Ibrahim et al. ("
+2003) would be hardly accounted by the fall-back model unless the accretion rate were decreasing very rapidly.,2003) would be hardly accounted by the fall-back model unless the accretion rate were decreasing very rapidly.
+ Moreover. the above active phase is expected to occur just once in the AXP life (with a duration longer than the about one year observed in aand possibly:: Chatterjee et al.," Moreover, the above active phase is expected to occur just once in the AXP life (with a duration longer than the about one year observed in and possibly: Chatterjee et al."
+ 2000)., 2000).
+ The pronounced long-term X-ray flux and pulsed fraction variability of mmight be more easily explained in the framework of a pulsar in a binary system with a light companion., The pronounced long-term X-ray flux and pulsed fraction variability of might be more easily explained in the framework of a pulsar in a binary system with a light companion.
+ In fact. we note that the quiescent luminosity of ((and )) is similar to that already observed from transient binary system pulsars (Campana et al.," In fact, we note that the quiescent luminosity of (and ) is similar to that already observed from transient binary system pulsars (Campana et al."
+ 2002). together with the pulsed fraction decrease in. the on/off transitions (Campana et al.," 2002), together with the pulsed fraction decrease in the on/off transitions (Campana et al."
+ 2001)., 2001).
+ However. in. the binary scenario. the short X-ray bursts displayed by AXPs and. especially. SGRs would be difficult to explain (see Mereghetti et al.," However, in the binary scenario, the short X–ray bursts displayed by AXPs and, especially, SGRs would be difficult to explain (see Mereghetti et al."
+ 2002 for more details)., 2002 for more details).
+ We thank Harvey Tananbaum for making the HHRC-I observation possible through the DDirector’s Discretionary Time program. and Michael Juda for help during the sscheduling phase.," We thank Harvey Tananbaum for making the HRC–I observation possible through the Director's Discretionary Time program, and Michael Juda for help during the scheduling phase."
+ We also thank the ESO Director's Discretionary Time Committee for accepting to extend our original program allowing us to trigger the VLT observations. and Lowell Taceconi-Garman for his help and patience during the OB submission phase.," We also thank the ESO Director's Discretionary Time Committee for accepting to extend our original program allowing us to trigger the VLT observations, and Lowell Tacconi-Garman for his help and patience during the OB submission phase."
+ This work is partially supported through ASI and Ministero dellUniversità e Ricerca Scientifica e Tecnologica (MURST-COFIN) grants., This work is partially supported through ASI and Ministero dell'Università e Ricerca Scientifica e Tecnologica (MURST–COFIN) grants.
+Alulticolour photometry can be decisive in solving this classification problem because Cepheids have characteristic amplitude ratios.,Multicolour photometry can be decisive in solving this classification problem because Cepheids have characteristic amplitude ratios.
+ Our new photometric data suggest. that V2210CC is not a Cepheid., Our new photometric data suggest that Cyg is not a Cepheid.
+ Phe amplitude. of. its rightness variation$8 in Vis only slightly smaller than the amplitude in the D band: the ratio is about 0.9 instead. of he usual value of 0.650.70 (Ixlagvivik&Szabacos2009)., The amplitude of its brightness variation in $V$ is only slightly smaller than the amplitude in the $B$ band: the ratio is about 0.9 instead of the usual value of $0.65-0.70$ \citep{klsz09}.
+". Although a bright blue companion star is able to suppress he observable amplitude of the light variations in the D ind. the observed amplitude ratio is incompatible with the Cepheid nature even if the star had à very hot companion. Roy. Hs,"," Although a bright blue companion star is able to suppress the observable amplitude of the light variations in the $B$ band, the observed amplitude ratio is incompatible with the Cepheid nature even if the star had a very hot companion. $R_{21}$,"
+" and ós, place it among the first overtone objects. out Os, ds very dillerent. from. both fundamental and. first overtone progressions reffps))."," $R_{31}$ and $\phi_{31}$ place it among the first overtone objects, but $\phi_{21}$ is very different from both fundamental and first overtone progressions \\ref{fps}) )."
+ In a next slop we analyze the evele-to-cvele variations in the periocicity using the so-called ο6 diagram., In a next step we analyze the cycle-to-cycle variations in the periodicity using the so-called $O-C$ diagram.
+ The O C' diagram constructed for the moments of brightness maxima is plotted in the top panel of retligv2279ockepler.., The $O-C$ diagram constructed for the moments of brightness maxima is plotted in the top panel of \\ref{Figv2279ockepler}.
+ Phe zero epoch was arbitrarily chosen at the first maximum in theAecpéer data set., The zero epoch was arbitrarily chosen at the first maximum in the data set.
+ X Least squares fit to the O—€C' residuals resulted. in the best. fitting period of 4.125642dd. which is somewhat longer than any formerly published value for V2279CCve.," A least squares fit to the $O-C$ residuals resulted in the best fitting period of d, which is somewhat longer than any formerly published value for Cyg."
+ We then included all eround-basecl observations we could find to construct a new O6 diagram using the ephemeris C=BID(2454954.6109+0.0023)|(4.125642£0.0000514.," We then included all ground-based observations we could find to construct a new $O-C$ diagram using the ephemeris $C = {\rm BJD}\,(2\,454\,954.6109 \pm 0.0023) + (4.125642 \pm 0.000051)
+E$."
+ refTabOCy2279 contains the moments of maxima. the corresponding epochs and the οC' residuals. of the available observations.," \\ref{TabOCv2279} contains the moments of maxima, the corresponding epochs and the $O-C$ residuals of the available observations."
+ We assigned. dilferent weights to data from cillerent sources., We assigned different weights to data from different sources.
+ These weights (between 1 and 5) correspond. to the ‘goodness’ of the seasonal light. curve. largerD number means better coverage5 and smaller scatter (forNepler data W=5).," These weights (between 1 and 5) correspond to the `goodness' of the seasonal light curve, larger number means better coverage and smaller scatter (for data W=5)."
+ The period of V2279CCve show large Uuctuations during the [ast decade (see the lower panel of refligv22790ckepler))., The period of Cyg show large fluctuations during the last decade (see the lower panel of \\ref{Figv2279ockepler}) ).
+ Our additional grouncd-basec photometric cata indicate| a —very recent —change in the period that we will investigate using future quarters of data., Our additional ground-based photometric data indicate a very recent change in the period that we will investigate using future quarters of data.
+ We note that the star has a high contamination index reftabl)). indicating that the ~1LS percent of the measures Ilux comes form other sources.," We note that the star has a high contamination index \\ref{tab1}) ), indicating that the $\sim$ 18 percent of the measured flux comes form other sources."
+ The spectrum of V2279CCve. olfers. the. cleares evidence that this star has been misclassified as a Cepheid., The spectrum of Cyg offers the clearest evidence that this star has been misclassified as a Cepheid.
+ We have plotted three segments of the spectrum taken with the spectrograph in ro[fv2270spec containing these lines: LULL ΑΛ). DD (5890 and AA)) and Li AA).," We have plotted three segments of the spectrum taken with the spectrograph in \\ref{v2279spec} containing these lines: H ), D (5890 and ) and Li )."
+ While the DD lines are normal. the characteristics of the two other lines do not support a Cepheic nature of CCve.," While the D lines are normal, the characteristics of the two other lines do not support a Cepheid nature of Cyg."
+ The Li AA)) is never seen inemission in Cepheids. while the Ca emission implies chromospheric activity. which is also not a Ceopheid characteristic.," The Li ) is never seen inemission in Cepheids, while the Ca emission implies chromospheric activity, which is also not a Cepheid characteristic."
+ We fitted. theoretical template spectra from the extensive spectral library of Alunarietal.(2005) to the spectrum ancl determined the following atmospheric parameters: Yigg=4900+ 200KW. leg= 37404. M/H]=12404 and esin;=40kkmss," We fitted theoretical template spectra from the extensive spectral library of \citet{mun05} to the spectrum and determined the following atmospheric parameters: $T_{\rm eff}=4900\pm200$ K, $\log{g}=3.7\pm0.4$ , $[M/H]=-1.2\pm0.4$ and $v \sin i = 40$ $^{-1}$."
+ The resulting. parameters are also incompatible with a Cepheid variable. but suggest a cool main-sequence star with moderate rotation.," The resulting parameters are also incompatible with a Cepheid variable, but suggest a cool main-sequence star with moderate rotation."
+ Summarizing the previous subsections. we conclude that all the candidates turned out not to be Cepheids except the already known CepheidV1I54C€€ve. which we will deseribe in detail in the following.," Summarizing the previous subsections, we conclude that all the candidates turned out not to be Cepheids except the already known CepheidCyg, which we will describe in detail in the following."
+" In this section we analyze both Acpler data and ground- follow-up observations of what is apparently the only Copheid being observed in FOV,", In this section we analyze both data and ground-based follow-up observations of what is apparently the only Cepheid being observed in FOV.
+analysed by the DS test following the same prescription as the real ones.,analysed by the DS test following the same prescription as the real ones.
+ In the simulated MC clusters we know whether the galaxies selected as substructure by the DS test are real or spurious substructure galaxies., In the simulated MC clusters we know whether the galaxies selected as substructure by the DS test are real or spurious substructure galaxies.
+ It is expected that the number of galaxies in substructure will strongly depend on the adopted value of 0..., It is expected that the number of galaxies in substructure will strongly depend on the adopted value of $\delta_{c}$.
+" The best 6,. values should maximize the detection of galaxies in substructure and minimize the spurious ones.", The best $\delta_{c}$ values should maximize the detection of galaxies in substructure and minimize the spurious ones.
+" Figure 3. shows the number of galaxies detected as substructure versus the number of spurious substructure detections for the simulated galaxy clusters and for the different values of 0,.", Figure \ref{f2} shows the number of galaxies detected as substructure versus the number of spurious substructure detections for the simulated galaxy clusters and for the different values of $\delta_{c}$.
+ The results from simulated clusters with 50 and 100 galaxies containing 10% or 30% of their population in substructures are shown in Fig. 3.., The results from simulated clusters with 50 and 100 galaxies containing $\%$ or $\%$ of their population in substructures are shown in Fig. \ref{f2}.
+ Similarly to what was reported in Sec., Similarly to what was reported in Sec.
+ 3.1. the DS test does not detect individual galaxies in substructures when the relaxed and the unrelaxed component have similar mean velocities.," 3.1, the DS test does not detect individual galaxies in substructures when the relaxed and the unrelaxed component have similar mean velocities."
+ It is also evident from Fig., It is also evident from Fig.
+" 3. that studies of substructure adopting a unique value of o. for all clusters would be almost dominated by spurious detections 1f 0,94 was adopted."," \ref{f2} that studies of substructure adopting a unique value of $\delta_{c}$ for all clusters would be almost dominated by spurious detections if $\delta_{g,95}$ was adopted."
+" The spurious detections are lower if 0,59 1s used."," The spurious detections are lower if $\delta_{g,99}$ is used."
+" Studies of substructure using individual values of 6, for each cluster will be almost free of spurious substructure galaxies by adopting Ojme.."," Studies of substructure using individual values of $\delta_{c}$ for each cluster will be almost free of spurious substructure galaxies by adopting $\delta_{i,max}$."
+ Nevertheless. the fraction of galaxies in substructure detected will be very low.," Nevertheless, the fraction of galaxies in substructure detected will be very low."
+ The mean fraction of galaxies in substructure detected using 0;oo is 14% higher than that adopted with 9;ooo.," The mean fraction of galaxies in substructure detected using $\delta_{i,99}$ is $\%$ higher than that adopted with $\delta_{i,99.9}$."
+ In contrast. the spurious detections are only 6% higher with 9;o9 than with 0;09.," In contrast, the spurious detections are only $\%$ higher with $\delta_{i,99}$ than with $\delta_{i,99.9}$."
+" Based on the previous considerations. we adopt 9;o9 and 6,9 for the determination of the galaxies in substructure."," Based on the previous considerations, we adopt $\delta_{i,99}$ and $\delta_{g,99}$ for the determination of the galaxies in substructure."
+ We computed the fraction of galaxies in substructure as a function of their location in the cluster., We computed the fraction of galaxies in substructure as a function of their location in the cluster.
+ Figure 4. shows the fraction of galaxies located in substructures for EC] and EC2 ensemble clusters., Figure \ref{f3} shows the fraction of galaxies located in substructures for EC1 and EC2 ensemble clusters.
+" In both ensemble clusters. the fraction of galaxies in substructure depends on the 9, value adopted."," In both ensemble clusters, the fraction of galaxies in substructure depends on the $\delta_{c}$ value adopted."
+ There is a clear dependence of the fraction of substructure and the location of the galaxies in the cluster., There is a clear dependence of the fraction of substructure and the location of the galaxies in the cluster.
+ In both ensemble clusters. the fraction of substructure in the inner cluster regions rooo) is smaller than in the regions located at ¢>7209.," In both ensemble clusters, the fraction of substructure in the inner cluster regions $r<r_{200}$ ) is smaller than in the regions located at $r>r_{200}$."
+ The EC2 ensemble cluster shows. independently of the ὃς value adopted. a higher fraction of substructure (especially in the outer cluster region).," The EC2 ensemble cluster shows, independently of the $\delta_{c}$ value adopted, a higher fraction of substructure (especially in the outer cluster region)."
+ This shows that the fraction of substructure grows as funt galaxies are included in the samples., This shows that the fraction of substructure grows as faint galaxies are included in the samples.
+ Note that we have not observed any segregation between galaxies in substructures and magnitudes (see Fig. 5))., Note that we have not observed any segregation between galaxies in substructures and magnitudes (see Fig. \ref{f4}) ).
+ This means that the fraction of faint galaxies in substructures is not higher than the fraction of bright ones., This means that the fraction of faint galaxies in substructures is not higher than the fraction of bright ones.
+ We have investigated the dependence of the previous results with the number of galaxies per cluster., We have investigated the dependence of the previous results with the number of galaxies per cluster.
+ In particular. we have," In particular, we have"
+column density field PDP than in the true density Field PDE.,column density field PDF than in the true density field PDF.
+ This suggests that observed positive tails in column density DES imply the presence of similar positive tails in density DES. but that some caution in their interpretation should »e applied.," This suggests that observed positive tails in column density PDFs imply the presence of similar positive tails in density PDFs, but that some caution in their interpretation should be applied."
+ The overall success of the method. supports the idea hat the form of the density PDE is. preserved. in. the column density PDE (even for a non-Iognormal form) with appropriate provisos on the extreme positive ancl negative ails cliscussed above., The overall success of the method supports the idea that the form of the density PDF is preserved in the column density PDF (even for a non-lognormal form) with appropriate provisos on the extreme positive and negative tails discussed above.
+ We have introduced. and tested. a simple. method for reconstructing the probability density function (PDE) of a 3D turbulent density field using information present solely in the projected (observable) column density field. in. 2D. The method: builds on a previously established: method to calculate the 3D normalised. density variance. recently presented by Brunt. Federrath. and Price (BET 2010)," We have introduced and tested a simple method for reconstructing the probability density function (PDF) of a 3D turbulent density field using information present solely in the projected (observable) column density field in 2D. The method builds on a previously established method to calculate the 3D normalised density variance, recently presented by Brunt, Federrath, and Price (BFP, 2010)."
+ To a good approximation. the PDE of log(normalised column density) is a compressed. shifted version of the PDL of log(normatisecl density). but can deviate significantly in the extreme tails.," To a good approximation, the PDF of log(normalised column density) is a compressed, shifted version of the PDF of log(normalised density), but can deviate significantly in the extreme tails."
+ The compression factor.ἕ can be derived observationally from the column density power spectrum. assuming statistical isotropy. using the BET method.," The compression factor, $\xi$, can be derived observationally from the column density power spectrum, assuming statistical isotropy, using the BFP method."
+ This work was supported by STEC Grant. ST/F003277/1., This work was supported by STFC Grant ST/F003277/1.
+ Wed like to thank the anonymous referee. for &ood suggestions that improved. the paper., We'd like to thank the anonymous referee for good suggestions that improved the paper.
+ CB is supported. by an RCUW fellowship at the University. of Lxcter. Ulx.," CB is supported by an RCUK fellowship at the University of Exeter, UK."
+ Ck is grateful for financial support by the International Alax Planck Research School for Astronomy and. Cosmic Physics (LAIPRS-A) and the Heidelberg Graduate School of Fundamental Physics (LIGSEDP). which is funded by the Execllenee Initiative of the German Research Foundation (DEG GSC 1290/1).," CF is grateful for financial support by the International Max Planck Research School for Astronomy and Cosmic Physics (IMPRS-A) and the Heidelberg Graduate School of Fundamental Physics (HGSFP), which is funded by the Excellence Initiative of the German Research Foundation (DFG GSC 129/1)."
+ The AHID simulations were run at the Leihniz-Reehenzentrum (grant pr32lo)., The MHD simulations were run at the Leibniz-Rechenzentrum (grant pr32lo).
+ The software used in this work was in part developed. by the DOL-supported ASC / Alliance Center for Astrophysical VPhermonuclear Flashes at the University of Chicago., The software used in this work was in part developed by the DOE-supported ASC / Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago.
+as the pressure-limiting expansion.,as the pressure-limiting expansion.
+ The argument for the pressure-limiting expansion can be understood as follows., The argument for the pressure-limiting expansion can be understood as follows.
+ Assume that the lobe expansion starts with near pressure balance p;~pex., Assume that the lobe expansion starts with near pressure balance $p_l\sim p_{\rm ex}$.
+" If a significant imbalance pi>pex develops as a result of substantial drop in the external pressure as the lobe expands into the underdense region, the expansion will accelerate and the system will relax to near pressure balance on a time much shorter than the age of the source."," If a significant imbalance $p_l>p_{\rm ex}$ develops as a result of substantial drop in the external pressure as the lobe expands into the underdense region, the expansion will accelerate and the system will relax to near pressure balance on a time much shorter than the age of the source."
+ One can derive the relevant expansion law from the energy conservation (Bicknell1986;Eilek&Shore 1989).," One can derive the relevant expansion law from the energy conservation \citep{b86,es89}."
+". Denoting the total energyof the lobe by Ej, the energy equation is given by where one ignores the radiative losses and assumes that the injection power is "," Denoting the total energyof the lobe by $E_l$, the energy equation is given by where one ignores the radiative losses and assumes that the injection power is $Q_j$."
+"It should be noted that p; is sum of all components Q;.(radiating particles, nonradiating particles, magnetic field)."," It should be noted that $p_l$ is sum of all components (radiating particles, nonradiating particles, magnetic field)."
+ Eq (3)) implies that the main cause for decrease in the internal energy in the lobe is the volume work done against the external pressure., Eq \ref{eq:El}) ) implies that the main cause for decrease in the internal energy in the lobe is the volume work done against the external pressure.
+" Assuming that the temperature of the external medium is constant, the external pressure can be written as Dex~ "," Assuming that the temperature of the external medium is constant, the external pressure can be written as $p_{\rm ex}\sim p_c(r_c/r)^{\beta}$."
+"One may writethe external pressure at the core as pe(re/r)?.p,=nokpTo©1.4x10711Pa for no= and Το=107K."," One may writethe external pressure at the core as $p_c=n_0k_BT_0\approx 1.4\times10^{-11}\,{\rm Pa}$ for $n_0=\rho_c/m_p=10^5\,{\rm m}^{-3}$ and $T_0=10^7\,\rm K$."
+" As we are interested in the pressure-limiting expansion with p;~ pex, one cansubstitute py~pe(r«/r)? for (3)) to obtain with In deriving (4)), we assume E;=p,V/(T—1) and V=xr?, where I is the adiabatic index of the plasma in the lobe."," As we are interested in the pressure-limiting expansion with $p_l\sim p_{\rm ex}$ , one cansubstitute $p_l\sim p_c(r_c/r)^{\beta}$ for \ref{eq:El}) ) to obtain with In deriving \ref{eq:Dt3}) ), we assume $E_l=p_lV/(\Gamma-1)$ and $V=\chi r^3$, where $\Gamma$ is the adiabatic index of the plasma in the lobe."
+ The geometry of the lobe is described by the parameter χ., The geometry of the lobe is described by the parameter $\chi$.
+" Since there is no strong observational evidence for the dependence of the aspect ratio on the size, it is justified to treat x as an independent parameter."," Since there is no strong observational evidence for the dependence of the aspect ratio on the size, it is justified to treat $\chi$ as an independent parameter."
+" As x=47/3 corresponds to a sphere, one generally has x«47/3."," As $\chi=4\pi/3$ corresponds to a sphere, one generally has $\chi\ll4\pi/3$."
+ A special case is the constant-pressure expansion (Eilek&Shore1989;KaiserBest2007).," A special case is the constant-pressure expansion \citep{es89,kb07}."
+". With pe,=const, one obtains Eq (6)) implies that the size of the radio sources driven(6) by low-power jets grows much more slowly than the sources driven by high-power jets."," With $p_{ex}=$ const, one obtains Eq \ref{eq:Dt4}) ) implies that the size of the radio sources driven by low-power jets grows much more slowly than the sources driven by high-power jets."
+" The total radio power as a function of time can be written as an integration of the single-particle power, P,, over the spatially-integrated, time-dependent particle distribution, N(y,t), where »y is the particle’s Lorentz factor."," The total radio power as a function of time can be written as an integration of the single-particle power, $P_s$, over the spatially-integrated, time-dependent particle distribution, $N(\gamma,t)$, where $\gamma$ is the particle's Lorentz factor."
+" In practice, one may assumes that each particle emits at the characteristic frequency v.=(3/4)vBy’, producing a power spectrum P&ó(v—νο), where vg is the gyrofrequency and P4 is the single-particle synchrotron power averaged on the particle's pitch angle, given by where στ is the Thomson cross section and Ug is the magnetic energy density."," In practice, one may assumes that each particle emits at the characteristic frequency $\nu_c=(3/4)\nu_B\gamma^2$, producing a power spectrum $P_s\delta(\nu-\nu_c)$, where $\nu_B$ is the gyrofrequency and $P_s$ is the single-particle synchrotron power averaged on the particle's pitch angle, given by where $\sigma_T$ is the Thomson cross section and $U_B$ is the magnetic energy density."
+" If the pitch angle distribution is maintained in the isotropic state throughout the evolution, the total spectral power can be expressed as where y,=(4v/3vg)!/? is the Lorentz factor of particles that emit synchrotron radiation predominantly at frequency v."," If the pitch angle distribution is maintained in the isotropic state throughout the evolution, the total spectral power can be expressed as where $\gamma_*\equiv(4\nu/3\nu_B)^{1/2}$ is the Lorentz factor of particles that emit synchrotron radiation predominantly at frequency $\nu$."
+" Eq can be used to eliminate ¢ in (8)), giving rise to P, as a (2))function of D."," Eq \ref{eq:Dt2}) ) can be used to eliminate $t$ in \ref{eq:Pnu}) ), giving rise to $P_\nu$ as a function of $D$."
+" All the variables in (8)) are global, in the sense that they can be regarded as spatial averages."," All the variables in \ref{eq:Pnu}) ) are global, in the sense that they can be regarded as spatial averages."
+ One can show that the synchrotron output (in radio) only comprises a tiny fraction of the total power input by the jet., One can show that the synchrotron output (in radio) only comprises a tiny fraction of the total power input by the jet.
+" The total number of relativistic particles that emit synchrotron radiation at frequency v is ~7Q;tyx”, where 7<1 is the fraction of the jet power into the relativistic particles, p is the particle spectral index and t is the age of the source."," The total number of relativistic particles that emit synchrotron radiation at frequency $\nu$ is $\sim \eta Q_j t\gamma^{-p}_*$, where $\eta\leq1$ is the fraction of the jet power into the relativistic particles, $p$ is the particle spectral index and $t$ is the age of the source."
+" Since the total synchrotron power is ~P,v, the ratio of the synchrotron power to the jet input (7Q;) is estimated as P,v/qQ;j~ "," Since the total synchrotron power is $\sim P_\nu\nu$, the ratio of the synchrotron power to the jet input $\eta Q_j$ ) is estimated as $P_\nu\nu/\eta Q_j\sim \sigma_TcU_Bt(\nu/\nu_B)^{1/2}$ ."
+"For B= 50nT, t= 1Myr,and v=1 GHz, this orvCUpt(v/vg)!?.ratio is 1074, which implies that only a tiny fraction of the jet power is converted to synchrotron radiation."," For $B=50\,\rm nT$ , $t=1\,\rm Myr$ ,and $\nu=1\,\rm GHz$ , this ratio is$\sim 10^{-4}$ , which implies that only a tiny fraction of the jet power is converted to synchrotron radiation."
+ 'The instrument sensitivity places a lower-limit on the observable flux density., The instrument sensitivity places a lower-limit on the observable flux density.
+" As a result, the P,—D tracks have a cut-off at which the flux density is too low to be"," As a result, the $P_\nu$ $D$ tracks have a cut-off at which the flux density is too low to be"
+mass Interior to (hat radius.,mass interior to that radius.
+" Equally so. 05, is measured at the apocenter of its ry orbit and therefore smaller then the corresponding circular orbit."," Equally so, $v_{bar}$ is measured at the apocenter of its $x_1$ orbit and therefore smaller then the corresponding circular orbit."
+ Therefore wwould be overestimated [rom a value consistent with the mass distribution., Therefore would be overestimated from a value consistent with the mass distribution.
+ Conversely Lor bars aligned along the minor axis would be underestimated., Conversely for bars aligned along the minor axis would be underestimated.
+ We did not see such a trend in Figure 1., We did not see such a trend in Figure 1.
+ We should also note that the rotation curve is derived [rom a (wo dimensional velocity field which will tend to average oul this effect., We should also note that the rotation curve is derived from a two dimensional velocity field which will tend to average out this effect.
+ Figure 1 shows that galaxies which have more mass concentrated in their bulges have rounder bars., Figure 1 shows that galaxies which have more mass concentrated in their bulges have rounder bars.
+ However. (he existence of a correlation does not necessarily imply a causal relationship.," However, the existence of a correlation does not necessarily imply a causal relationship."
+ For example. it might be (hat some other parameter used to caleulate (he masses is more revelant.," For example, it might be that some other parameter used to calculate the masses is more revelant."
+" In particular. the central mass concentration is calculated as Pinc=ogHyg]UnsCB> it might be that the relative bulge size /.=Iyj/Tij, 1s more relevant than the dynamical mass concentration."," In particular, the central mass concentration is calculated as ${v_{blg}}^{2}R_{blg}/{v_{bar}}^{2}R_{bar}$; it might be that the relative bulge size $l_{c}=R_{blg}/R_{bar}$ is more relevant than the dynamical mass concentration."
+To investigate (his. we plotted ellipticities against /. for 25 barred SONG ealaxies (Figure 2).,"To investigate this, we plotted ellipticities against $l_{c}$ for 25 barred SONG galaxies (Figure 2)."
+" Comparison with Figure 1 shows that the correlation of ellipticity with /. is not as good as with f/,,:: the linear correlation coefficient. (nnweighted) is —0.56. significantly less than for/,,."," Comparison with Figure 1 shows that the correlation of ellipticity with $l_{c}$ is not as good as with $f_{mc}$ ; the linear correlation coefficient (unweighted) is $-$ 0.56, significantly less than for."
+.. We therefore conclude that. iis more likely than ἐς to be relevant for (he central mass concentration., We therefore conclude that is more likely than $l_{c}$ to be relevant for the central mass concentration.
+ The points in Figure 2 are coded to represent the 3 IIubble (vpes ie. early. intermediate and late (wpe barred ealaxies.," The points in Figure 2 are coded to represent the 3 Hubble types i.e. early, intermediate and late type barred galaxies."
+ Early (vpe galaxies have large. bright bulges whereas late (wpe galaxies appear to have less prominent bulges (ΕΠΙΡίο 1926).," Early type galaxies have large, bright bulges whereas late type galaxies appear to have less prominent bulges (Hubble 1926)."
+ In Figure 2. the different (vpes are dispersed over the whole plot: (here does not appear to be any correlation of bar ellipticity will Hubble type.," In Figure 2, the different types are dispersed over the whole plot; there does not appear to be any correlation of bar ellipticity with Hubble type."
+ This may because the IIubble sequence is based on the optical brightness of bulges which mav not always be a good measure of (he bulge mass., This may because the Hubble sequence is based on the optical brightness of bulges which may not always be a good measure of the bulge mass.
+ We thus conclude that bar ellipticity is less correlated with bulee size or IIubble tvpe than with central mass concentration., We thus conclude that bar ellipticity is less correlated with bulge size or Hubble type than with central mass concentration.
+ Thus Figure 1 indicates Chat bar ellipticity. which approximately measures bar strength. decreases with mass concentration in the bar.," Thus Figure 1 indicates that bar ellipticity, which approximately measures bar strength, decreases with mass concentration in the bar."
+ The correlation is perhaps not suprising. since a spherical mass concentration in the center of a bar will tend to decrease the non-axisvinimetric effect of the bar.," The correlation is perhaps not suprising, since a spherical mass concentration in the center of a bar will tend to decrease the non-axisymmetric effect of the bar."
+ Nonetheless. it is worth pointing out (hat this correlation is predicted by secular evolution models. as we now discuss.," Nonetheless, it is worth pointing out that this correlation is predicted by secular evolution models, as we now discuss."
+ Simulations show that bars drive gas inwards. resulting in central star formation and a consequent buildup of the central mass.," Simulations show that bars drive gas inwards, resulting in central star formation and a consequent buildup of the central mass."
+ The central mass concentration affects the stability of the bar supporting cj orbits so that the bar finally dissolves leading to(he transformation, The central mass concentration affects the stability of the bar supporting $x_1$ orbits so that the bar finally dissolves leading tothe transformation
+with MERLIN. and would set severe limits to the breadth of field in à real observation.,"with MERLIN, and would set severe limits to the breadth of field in a real observation."
+ In the present case. the field of interest Is narrow. so coarse values were chosen in order to reduce the computing load.," In the present case, the field of interest is narrow, so coarse values were chosen in order to reduce the computing load."
+ The chosen values were the largest ones consistent with accurate gridding of the relatively small area of the field occupied by the simulated sources., The chosen values were the largest ones consistent with accurate gridding of the relatively small area of the field occupied by the simulated sources.
+ The simulated visibilities were gridded with uniform weighting and subjected separately to 2000 cycles of Hóggbom cleaning at a loop gain of 0.1 versus the same number of cycles. at the same gain. of Sault-Wieringa cleaning. using ΕΤ basis functions as described in section 3.1.. the F functions being Chebyshev polynomials (although to the small order used. these are scarcely distinguishable from Taylor series terms) and the 7 functions being the half-frequency Fourier cosine functions of equation 7..," The simulated visibilities were gridded with uniform weighting and subjected separately to 2000 cycles of Höggbom cleaning at a loop gain of 0.1 versus the same number of cycles, at the same gain, of Sault-Wieringa cleaning, using $FT$ basis functions as described in section \ref{conway}, , the $F$ functions being Chebyshev polynomials (although to the small order used, these are scarcely distinguishable from Taylor series terms) and the $T$ functions being the half-frequency Fourier cosine functions of equation \ref{equ_m}."
+ A number of images relating to this simulation are shown in figure 6.., A number of images relating to this simulation are shown in figure \ref{fig_d}.
+ An RMS value was calculated for those pixels within a rectangular area which is shown on the plots by a black outline., An RMS value was calculated for those pixels within a rectangular area which is shown on the plots by a black outline.
+ These numbers are shown in table {.., These numbers are shown in table \ref{tab_a}.
+ The Sault-Wieringa technique clearly offers à significant improvement in dynamic range., The Sault-Wieringa technique clearly offers a significant improvement in dynamic range.
+ As discussed in section 3.4.. the value of this technique lies more in its ability to remove artifacts from an image of average flux density. rather than as à way to estimate light curves or spectra.," As discussed in section \ref{reconstruction}, the value of this technique lies more in its ability to remove artifacts from an image of average flux density, rather than as a way to estimate light curves or spectra."
+ For wide-band observations in which there is no time variation over the observation. because of the typically slowly-varying nature of radio spectra. one may reasonably expect to extract low-order spectral information (such as spectral indices) from Conway decomposition.," For wide-band observations in which there is no time variation over the observation, because of the typically slowly-varying nature of radio spectra, one may reasonably expect to extract low-order spectral information (such as spectral indices) from Conway decomposition."
+ But because light curves may easily contain significant power at higher orders of the time basis functions. which are not so accurately recovered by the deconvolution. more caution is advisable in interpreting the time-dependent output of the deconvolution process.," But because light curves may easily contain significant power at higher orders of the time basis functions, which are not so accurately recovered by the deconvolution, more caution is advisable in interpreting the time-dependent output of the deconvolution process."
+ If an accurate light curve ofa source is desired. it is probably always," If an accurate light curve ofa source is desired, it is probably always"
+Cygnus X-I is à 5.6-day X-ray binary that harbors the best studied black-hole candidate in the Galaxy.,Cygnus X-1 is a 5.6-day X-ray binary that harbors the best studied black-hole candidate in the Galaxy.
+ The mass of the acereting compact object has been estimated as ~10.1 .. and the donor star ts classified as an O9.7 [ab supergiant of ~17.8 M.. (Herrero et al., The mass of the accreting compact object has been estimated as $\sim 10.1$ $_{\odot}$ and the donor star is classified as an O9.7 Iab supergiant of $\sim 17.8$ $_{\odot}$ (Herrero et al.
+ 1995)., 1995).
+ The system is located at ~2 kpe (e.g. Gierlinsski et al., The system is located at $\sim 2$ kpc (e.g. Gierlińsski et al.
+ 1999)., 1999).
+ The black hole accretes through the wind of the companion star., The black hole accretes through the wind of the companion star.
+ Most of the time. the X-ray source is in the so-called low/hard state. characterized by a relatively weak blackbody component peaking at a few keV plus a strong hard power-law of photon index ~1.6.," Most of the time, the X-ray source is in the so-called low/hard state, characterized by a relatively weak blackbody component peaking at a few keV plus a strong hard power-law of photon index $\sim1.6$."
+ A nonthermal radio Jet has been observed in this state (Stirling et al., A nonthermal radio jet has been observed in this state (Stirling et al.
+ 2001). extending up to 15 mas.," 2001), extending up to $\sim 15$ mas."
+ The jet seems to form an average angle with the line of sight of ~30° (Fender 2001) and it has been suggested that it might be precessing (Stirling et al., The jet seems to form an average angle with the line of sight of $\sim 30^{\circ}$ (Fender 2001) and it has been suggested that it might be precessing (Stirling et al.
+ 2001. Romero et al.," 2001, Romero et al."
+ 2002)., 2002).
+ Occasionally. a transition to a high/soft state can occur.," Occasionally, a transition to a high/soft state can occur."
+ In this state. most of the radiated energy is concentrated in the blackbody. while the power-law component becomes softer. with an index of ~2.8 and no jet has been observed.," In this state, most of the radiated energy is concentrated in the blackbody, while the power-law component becomes softer, with an index of $\sim 2.8$ and no jet has been observed."
+ The usual interpretation of the X-ray behavior of the source is that the blackbody component originates m a cold. optically thick accretion disk. whereas the power-law component is produced in an optically thin hot corona by thermal Comptonization of disk photons (Poutanen et al.," The usual interpretation of the X-ray behavior of the source is that the blackbody component originates in a cold, optically thick accretion disk, whereas the power-law component is produced in an optically thin hot corona by thermal Comptonization of disk photons (Poutanen et al."
+ 1997. Dove et al.," 1997, Dove et al."
+ 1997. Esin et al.," 1997, Esin et al."
+ 1997. 1998).," 1997, 1998)."
+ The hot corona fills the inner few tens of gravitational radii around the compact object and the accretion disk would penetrate only marginally into the coronal region., The hot corona fills the inner few tens of gravitational radii around the compact object and the accretion disk would penetrate only marginally into the coronal region.
+ In the low/hard state the thermal X-ray luminosity is dominated by the corona. with typical luminosities of ~ a few times 10? erg s7!.," In the low/hard state the thermal X-ray luminosity is dominated by the corona, with typical luminosities of $\sim$ a few times $10^{37}$ erg $^{-1}$."
+ During the transition to the high/soft state the corona is likely ejected as the accretion disk approaches to the black hole (Fender et al., During the transition to the high/soft state the corona is likely ejected as the accretion disk approaches to the black hole (Fender et al.
+ 2004. also Esin et al.," 2004, also Esin et al."
+ 1997. 1998).," 1997, 1998)."
+ Most of the energy is then dissipated by the disk. until the inner part of it dominates the radiation again. and the cycle starts again.," Most of the energy is then dissipated by the disk, until the inner part of it dominates the radiation again, and the cycle starts again."
+ In the low/hard state. the disk is illuminated by hard photons from the corona resulting in the production of an Fe Kev line and a Compton reflection feature.," In the low/hard state, the disk is illuminated by hard photons from the corona resulting in the production of an Fe $\alpha$ line and a Compton reflection feature."
+ The first detection of the line was made by Barr et al. (, The first detection of the line was made by Barr et al. (
+1985) with EXOSAT.,1985) with $EXOSAT$.
+ They reported a broad (FWHM ~1.2 keV) emission line at 6.2 keV with an equivalent width of ~120 eV. Kitamoto et al. (, They reported a broad (FWHM $\sim1.2$ keV) emission line at $\sim 6.2$ keV with an equivalent width of $\sim120$ eV. Kitamoto et al. (
+1990) obtained a Tenma GSPC spectrum which was consistent with a narrow emission line at 6.5 keV with an equivalent width of 60—80 eV. A Compton reflection feature was then found above 20 keV (see Tanaka 199] and references therein).,1990) obtained a $Tenma$ GSPC spectrum which was consistent with a narrow emission line at $\sim 6.5$ keV with an equivalent width of $60-80$ eV. A Compton reflection feature was then found above 20 keV (see Tanaka 1991 and references therein).
+ The Ginga spectrum in the 2—30 keV range can be fitted quite well by the sum of a power-law with index. ~1.7. a reflection component and a narrow Fe emission line at 6.4 keV with an equivalent width of ~60 eV (Tanaka 1991).," The $Ginga$ spectrum in the $2-30$ keV range can be fitted quite well by the sum of a power-law with index $\sim1.7$, a reflection component and a narrow Fe emission line at $6.4$ keV with an equivalent width of $\sim 60$ eV (Tanaka 1991)."
+ Subsequent ASCA observations confirmed these results but restricting the width to 10—30 eV (Ebisawa et al., Subsequent $ASCA$ observations confirmed these results but restricting the width to $10-30$ eV (Ebisawa et al.
+ 1996)., 1996).
+ A broad edge at E>7 keV was also reported., A broad edge at $E>7$ keV was also reported.
+ À detailed historical account of attempts to detect an Fe line in Cygnus X-1 and the intrinsic difficulties it entails. can be found in Reynolds Nowak (2003).," A detailed historical account of attempts to detect an Fe line in Cygnus X-1 and the intrinsic difficulties it entails, can be found in Reynolds Nowak (2003)."
+ Recently. observed Cygnus X-I with the High Energy Transmission Grating Spectrometer in an intermediate X-ray state (Miller et al.," Recently, observed Cygnus X-1 with the High Energy Transmission Grating Spectrometer in an intermediate X-ray state (Miller et al."
+ 2002)., 2002).
+ The narrow Fe line was detected at £26.4153-0.007 keV with an equivalent width of W=1653 eV. along with a broad line at E=5.82+0.07 keV with W= eV. A smeared edge was also detected at 7.30.2 keV.1011 Miller et al. (," The narrow Fe line was detected at $E=6.415\pm
+0.007$ keV with an equivalent width of $W=16^{+3}_{-2}$ eV, along with a broad line at $E=5.82\pm 0.07$ keV with $W=140^{+70}_{-40}$ eV. A smeared edge was also detected at $7.3\pm 0.2$ keV. Miller et al. ("
+2002) interpret these results in terms of an accretion disk with irradiation of the inner disk producing the broad Fe Ko emission line and edge. and irradiation of the outer disk producing the narrow line.,"2002) interpret these results in terms of an accretion disk with irradiation of the inner disk producing the broad Fe $\alpha$ emission line and edge, and irradiation of the outer disk producing the narrow line."
+ The broad line is shaped by Doppler and gravitational effects and. to a lesser extent. by Compton reflection.," The broad line is shaped by Doppler and gravitational effects and, to a lesser extent, by Compton reflection."
+ Por different spectral states. the different disk structure may change the Fe Ka line.," For different spectral states, the different disk structure may change the Fe $\alpha$ line."
+ For example. if the disk is truncated at a much larger radius rather than the innermost stable orbit às suggested for the hard state. then the width of the Fe Ko line may significantly decrease.," For example, if the disk is truncated at a much larger radius rather than the innermost stable orbit as suggested for the hard state, then the width of the Fe $\alpha$ line may significantly decrease."
+ In any case. the variation due to disk precession. if that were observable. and that due to the aceretion mode (disk structure). will have a different temporal signature (periodicity). what would make them easy to distinguish.," In any case, the variation due to disk precession, if that were observable, and that due to the accretion mode (disk structure), will have a different temporal signature (periodicity), what would make them easy to distinguish."
+ In this paper. we show that changes in the orientation. of the inner accretion disk of Cygnus X-] would affect the shape of the broad Fe Ko emission line in a periodic way.," In this paper, we show that changes in the orientation of the inner accretion disk of Cygnus X-1 would affect the shape of the broad Fe $\alpha$ emission line in a periodic way."
+about fitting scaling relations with measurement errors and intrinsic scatter iu both quantities.,about fitting scaling relations with measurement errors and intrinsic scatter in both quantities.
+ AD eratefully acknowledges partial support οι INFN eraut. PD51., AD gratefully acknowledges partial support from INFN grant PD51.
+ We thank Susan Tokuz for reducing the spectroscopic data and Perry Berlind aud Mike. Calkius for assisting with the observations., We thank Susan Tokarz for reducing the spectroscopic data and Perry Berlind and Mike Calkins for assisting with the observations.
+Planetary nebulae (PNe) are known to possess a variety of small-scale structures that are usually in a lower ionization state than the main body of the nebulae.,Planetary nebulae (PNe) are known to possess a variety of small-scale structures that are usually in a lower ionization state than the main body of the nebulae.
+ Phe morphological and kinematic properties of these Iow-ionization structures (LISs. Goncaalves. Corradi Alampaso. 2001) vary. from tvpe to type. in the sense that LISs can appear in the form of pairs of knots. filaments. jets. and isolated features moving with velocities that either do not diller substantially from that of the ambient nebula. or instead. move supersonically through the environment.," The morphological and kinematic properties of these low-ionization structures (LISs, Gonçaalves, Corradi Mampaso, 2001) vary from type to type, in the sense that LISs can appear in the form of pairs of knots, filaments, jets, and isolated features moving with velocities that either do not differ substantially from that of the ambient nebula, or instead move supersonically through the environment."
+ The total number of PNe that are known to possess LISs is 55. Le. about of all the 527 Galactic PNe imaged in filters of high- and low-ionization emission lines (Balick LOST: Schwarz. Corradi Alelnick 1992: Manchado et al.," The total number of PNe that are known to possess LISs is 55, i.e. about of all the 527 Galactic PNe imaged in filters of high- and low-ionization emission lines (Balick 1987; Schwarz, Corradi Melnick 1992; Manchado et al."
+ 1996)., 1996).
+ The cülferent tvpes of LISs may be easily seen in Corradi et al. (, The different types of LISs may be easily seen in Corradi et al. (
+1996).,1996).
+ At present. the origin of jets ancl pairs of knots in PNe is not completely clear.," At present, the origin of jets and pairs of knots in PNe is not completely clear."
+ From the theoretical point. of view. the principal physical. process behind the formation of collimated LISs is the interplay between the stellar AGB and postACGD winds (for single stars) or between stellar aud disk winds (i£ the central star is binary).," From the theoretical point of view, the principal physical process behind the formation of collimated LISs is the interplay between the stellar AGB and post-AGB winds (for single stars) or between stellar and disk winds (if the central star is binary)."
+ According to the various studies dedicated to this issue (c.g. a-Segura 1997: a-segura Loppez 2000: Stellen. Lóppez Lim 2001: Blackman. Frank Welch 2001) jets ancl knots originated bv this interplay are. predicted. to. be supersonic. highly collimated and. two-sided.," According to the various studies dedicated to this issue (e.g. a-Segura 1997; a-Segura Lóppez 2000; Steffen, Lóppez Lim 2001; Blackman, Frank Welch 2001) jets and knots originated by this interplay are predicted to be supersonic, highly collimated and two-sided."
+ In the case of single stars these are expected to be produced at the same time as the main PN shell. but in the case of a binary star origin they may be vounger than the host PN.," In the case of single stars these are expected to be produced at the same time as the main PN shell, but in the case of a binary star origin they may be younger than the host PN."
+ However. important properties of these LISs such as density contrasts. the peculiar nitrogen abundance and main excitation mechanisms. appear hard to explain (see Dwarkadas Balick 1998: Dalick Frank 2002: CGoncaalves et al.," However, important properties of these LISs such as density contrasts, the peculiar nitrogen abundance and main excitation mechanisms, appear hard to explain (see Dwarkadas Balick 1998; Balick Frank 2002; Gonçaalves et al."
+ 2003)., 2003).
+ Α΄ 4-47 (DN 1100|04.3). is à compact PN that contains λος., K 4-47 (PN G149.0+04.4) is a compact PN that contains LISs.
+ It is composed. by a small. high. ionization nebular core and a pair of low-ionization. high-velocity," It is composed by a small, high ionization nebular core and a pair of low-ionization, high-velocity"
+times lower than for matching clusters without it.,times lower than for matching clusters without it.
+ Except for the largest cluster where the massive central galaxy produces a deep potential well the X-ray luminosity profile is less centrally concentrated than in the non-cooling case with a greater contribution coming from larger raclii., Except for the largest cluster where the massive central galaxy produces a deep potential well the X-ray luminosity profile is less centrally concentrated than in the non-cooling case with a greater contribution coming from larger radii.
+ Vhis cllect assists in convergence as we are less dependent upon the very centre of the cluster profile. (, This effect assists in convergence as we are less dependent upon the very centre of the cluster profile. (
+b) The spread of the X-ray luminosity temperature relation is well reproduced by our clusters.,b) The spread of the X-ray luminosity – temperature relation is well reproduced by our clusters.
+ Our non-cooling clusters lie close to the regression line suggested by 2] and have a similar slope (pxor 7)., Our non-cooling clusters lie close to the regression line suggested by \cite{E98} and have a similar slope $\rho \propto r^{-2}$ ).
+ We suggest that the increasing dominance of a large central galaxy on the local potential may produce the luminosity excess that drives the observed. X-ray. luminosity temperature relation away from the theoretically predicted. slope., We suggest that the increasing dominance of a large central galaxy on the local potential may produce the luminosity excess that drives the observed X-ray luminosity – temperature relation away from the theoretically predicted slope.
+Microscopie predictions of (he nuclear equation of state (Το). together with empirical constraints Irom Eos-sensilive observables. are a powerful combination to learn about the in-imecdium behavior of the nuclear force.,"Microscopic predictions of the nuclear equation of state (EoS), together with empirical constraints from EoS-sensitive observables, are a powerful combination to learn about the in-medium behavior of the nuclear force."
+ With this objective in mind. over the past several vears. our eroup has taken a broad look at the Eos exploring diverse aspects and extreme states of nuclear matter.," With this objective in mind, over the past several years, our group has taken a broad look at the EoS exploring diverse aspects and extreme states of nuclear matter."
+ From the experimental side. intense effort is going on to obtain reliable empirical information for the less known aspects of the EoS. Ileavy-ijon. (III) reactions are a popular wav lo seek constraints on (he svimmetry energy. through analvses of observables that are sensitive to (he pressure gradient between nuclear and neutron malter.," From the experimental side, intense effort is going on to obtain reliable empirical information for the less known aspects of the EoS. Heavy-ion (HI) reactions are a popular way to seek constraints on the symmetry energy, through analyses of observables that are sensitive to the pressure gradient between nuclear and neutron matter."
+ Isospin diffusion data from III collisions. together with analvses based on isospin-dependent transport models. provide information on the slope of Che symmetry energy.," Isospin diffusion data from HI collisions, together with analyses based on isospin-dependent transport models, provide information on the slope of the symmetry energy."
+ Naturally. different reaction conditions. in terms of energy per nucleon and/or impact parameter. will probe different densitv regions.," Naturally, different reaction conditions, in terms of energy per nucleon and/or impact parameter, will probe different density regions."
+ Concerning the lower densities. isospin-sensitive observables can also be identilied among the properties of normal nuclei.," Concerning the lower densities, isospin-sensitive observables can also be identified among the properties of normal nuclei."
+" The neutron skin of neutron-rich nuclei is a powerful isovector observable. being sensitive to the slope of the symmetry. energy. which determines to which extent neutrons are ""pushed out to form (the skin."," The neutron skin of neutron-rich nuclei is a powerful isovector observable, being sensitive to the slope of the symmetry energy, which determines to which extent neutrons are “pushed out"" to form the skin."
+ It is the purpose of (his note to svstematically examine and discuss the symmetry energy properties in microscopic models and the corresponding neutron skin predictions., It is the purpose of this note to systematically examine and discuss the symmetry energy properties in microscopic models and the corresponding neutron skin predictions.
+" We will take the skin of ΤΟ ΡΕ as our representative isovector ""observable""."," We will take the skin of $^{208}$ Pb as our representative isovector “observable""."
+ Paritv-violating electron scattering experiments are now a realistic option to determine neutron distributions with unprecedented accuracy., Parity-violating electron scattering experiments are now a realistic option to determine neutron distributions with unprecedented accuracy.
+ The neutron radius of 7 pb is expected to be measured within 0.05 [m thanks to (he electroweak program al the, The neutron radius of $^{208}$ Pb is expected to be measured within 0.05 fm thanks to the electroweak program at the
+2h* Mpe for anv range of halo masses (Tasitsioniietal.2004).,$\sim 2 h^{-1}$ Mpc for any range of halo masses \citep{taz:dissipationless}.
+ In this case. measurement ol large-scale galaxv-galaxy lensing and of the galaxy autocorrelation function determine the mass correlation [unction. O7Cos=(Qu£s)/£gg and so fixing the of the linear correlation function determines (he amplitude or normalizüion Όσες.," In this case, measurement of large-scale galaxy-galaxy lensing and of the galaxy autocorrelation function determine the mass correlation function, $\Omega_m^2 \xi_{mm} = (\Omega_m
+\xi_{gm})^2/\xi_{gg}$ and so fixing the of the linear correlation function determines the amplitude or normaliztion $\Omega_m\sigma_8$."
+" It also determines the bias. b(r)/Qun=£,,/0,,£,,."," It also determines the bias, $b(r)/\Omega_m =
+\xi_{gg}/\Omega_m\xi_{gm}$."
+ The assumption that bias is scale independent on large scales is crucial to extracting cosmological information from the galaxy. power spectrum: lensing can provide a wav to test this important. assumption., The assumption that bias is scale independent on large scales is crucial to extracting cosmological information from the galaxy power spectrum; lensing can provide a way to test this important assumption.
+ Other cosmological probes. such as cluster counts and cosmic shear. measure a different parameter combination. Oo. so combining these probes with cross-correlation lensing helps break this classic degeneracy.," Other cosmological probes, such as cluster counts and cosmic shear, measure a different parameter combination, $\Omega_m^{0.5} \sigma_8$, so combining these probes with cross-correlation lensing helps break this classic degeneracy."
+ An advantage of this new method is that it is robust: the only assumptions used are that general relativity is correct and that rí«(r)£z1 on large scales., An advantage of this new method is that it is robust; the only assumptions used are that general relativity is correct and that $\mbox{r}_{\times}(r)\simeq 1$ on large scales.
+ Applied to galaxy clusters. this inversion method should prove useful for both astrophysics and cosmology.," Applied to galaxy clusters, this inversion method should prove useful for both astrophysics and cosmology."
+ On small scales. cross-correlation lensing probes the mean density profiles of clusters.," On small scales, cross-correlation lensing probes the mean density profiles of clusters."
+ Cosmological dark matter simulations indicate (hat clusters have universal clark malter halos (hough barvonic physics presumably needs (o be taken into account to understand (he inner density structure in detail)., Cosmological dark matter simulations indicate that clusters have universal dark matter halos (though baryonic physics presumably needs to be taken into account to understand the inner density structure in detail).
+" Inverted density profiles can test this assumption quite directly,", Inverted density profiles can test this assumption quite directly.
+ As we have demonstrated. cluster mass profiles can be determined. leading to virial mass estimates independent of a model for the density. profiles.," As we have demonstrated, cluster mass profiles can be determined, leading to virial mass estimates independent of a model for the density profiles."
+ Virial masses allow a clirect comparison between simulated and real clusters., Virial masses allow a direct comparison between simulated and real clusters.
+ In addition. in a cluster survey. we can measure (he cluster abundance. n6M.z). binned by any observable proxy for mass. and use (his inversion method (o calibrate (he virial mass-observable relation as a function of redshift.," In addition, in a cluster survey we can measure the cluster abundance, $n(M,z)$, binned by any observable proxy for mass, and use this inversion method to calibrate the virial mass-observable relation as a function of redshift."
+ This should allow cluster survevs to more precisely probe cosmology. including the dark energy.," This should allow cluster surveys to more precisely probe cosmology, including the dark energy."
+ We still need to gather information about the scatter in the mass-observable relation: this can be constrained by simulations. bv self-calibration (Lima&Ili2005).. and internally by dividing (he clusters into subsamples and separately estimating (he relation.," We still need to gather information about the scatter in the mass-observable relation; this can be constrained by simulations, by self-calibration \citep{limahu2}, and internally by dividing the clusters into subsamples and separately estimating the mass-observable relation."
+ Belore applving (his method with confidence to real cluster data. the possible selection biases involved in a given cluster algorithm need to be explored.," Before applying this method with confidence to real cluster data, the possible selection biases involved in a given cluster algorithm need to be explored."
+ If the selection algorithm: prelerentially finds clusters aligned. along (he line of sight. the assumption of statistical isotropy will be violated ancl (he mass inversions mav be biased.," If the selection algorithm preferentially finds clusters aligned along the line of sight, the assumption of statistical isotropy will be violated and the mass inversions may be biased."
+ Similarly. il the selection algorithm assigns cluster centroids that are displaced from the true halo centers. the resulting density profiles will be convolved with distribution of centroid errors.," Similarly, if the selection algorithm assigns cluster centroids that are displaced from the true halo centers, the resulting density profiles will be convolved with distribution of centroid errors."
+ Additional complications for the cluster counting technique arise if clusters. however selected. are not isomorphic to massive dark matter halos.," Additional complications for the cluster counting technique arise if clusters, however selected, are not isomorphic to massive dark matter halos."
+ These issues can be tested with simulations., These issues can be tested with simulations.
+ Applving these lens inversion methods to clusters on large scales provides new cosmological information bevond the cluster counting technique., Applying these lens inversion methods to clusters on large scales provides new cosmological information beyond the cluster counting technique.
+ As noted lor galaxies above. we can use," As noted for galaxies above, we can use"
+be radially truncated in the low-hard state.,be radially truncated in the low–hard state.
+ Indeed. in some black holes observed at similar fractions of the Eddington limit (e.g.. XTE 15150. MeClintock et 22001). and at lower luminosities. 1t is possible that the disk may be radially recessed.," Indeed, in some black holes observed at similar fractions of the Eddington limit (e.g., XTE $+$ 480, McClintock et 2001), and at lower luminosities, it is possible that the disk may be radially recessed."
+ However. our results suggest that this prediction may not hold universally in the low—hard state.," However, our results suggest that this prediction may not hold universally in the low–hard state."
+" A phenomenological model for jet production (Fender. Belloni. Gallo 2004) suggests that a truncated disk may facilitate jet production. in part because Jet production does appear to be largely quenched in ""high-soft"" states of stellar-mass black holes (Fender 2005) wherein the inner disk is commonly thought to extend close to the black hole."," A phenomenological model for jet production (Fender, Belloni, Gallo 2004) suggests that a truncated disk may facilitate jet production, in part because jet production does appear to be largely quenched in “high–soft” states of stellar–mass black holes (Fender 2005) wherein the inner disk is commonly thought to extend close to the black hole."
+ Our observation of SWIFT J1753.5-0127 suggests that the jet production may not be enabled by a radially truncated disk. nor quenched by a filled innermost accretion disk.," Our observation of SWIFT $-$ 0127 suggests that the jet production may not be enabled by a radially truncated disk, nor quenched by a filled innermost accretion disk."
+ Other factors. perhaps including the mass accretion rate. or the absence of a strong corona in high-soft states. may inhibit jet production.," Other factors, perhaps including the mass accretion rate, or the absence of a strong corona in high–soft states, may inhibit jet production."
+ As previously noted. the spectral results obtained from this analysis are not necessarily at odds with inferences drawn from X-ray timing analyses of black holes in the low—hard state (Miller et 22006).," As previously noted, the spectral results obtained from this analysis are not necessarily at odds with inferences drawn from X-ray timing analyses of black holes in the low–hard state (Miller et 2006)."
+ Low-frequency QPOs and noise components are not likely to be directly related to Keplerian orbital frequencies at the ISCO., Low-frequency QPOs and noise components are not likely to be directly related to Keplerian orbital frequencies at the ISCO.
+ In transitions from the states to higher flux states. the frequency of low-frequency QPOs are sometimes observed to saturate.," In transitions from the states to higher flux states, the frequency of low-frequency QPOs are sometimes observed to saturate."
+ However. this is again not necessarily indicative of the innermost stable circular orbit. as such QPOs have been observed simultaneously with high-frequency QPOs which are plausibly related to Keplerian orbits in the inner disk (Remillard et 22002).," However, this is again not necessarily indicative of the innermost stable circular orbit, as such QPOs have been observed simultaneously with high-frequency QPOs which are plausibly related to Keplerian orbits in the inner disk (Remillard et 2002)."
+ On theoretical grounds. it is all but impossible that a standard optically-thick accretion disk can remain near the black hole in the least luminous phases of the low-hard state as the source approaches quiescence.," On theoretical grounds, it is all but impossible that a standard optically-thick accretion disk can remain near the black hole in the least luminous phases of the low–hard state as the source approaches quiescence."
+ However. the fraction of the Eddington lummosity at which the disk Is actually truncated is yet to be determined.," However, the fraction of the Eddington luminosity at which the disk is actually truncated is yet to be determined."
+ Exploring the regime near Ly/Legy107 will require an order of magnitude longer observation than this 43 ksee exposure., Exploring the regime near $L_X/L_{Edd} \simeq 10^{-4}$ will require an order of magnitude longer observation than this 43 ksec exposure.
+ Similarly. to obtain confirmation of the disk radius at LyΣΕ via the width of an Fe K emission line. would likely also require a 400—500 ksee observation.," Similarly, to obtain confirmation of the disk radius at $L_X/L_{Edd} \simeq 10^{-3}$ via the width of an Fe K emission line, would likely also require a 400–500 ksec observation."
+ We thank Norbert Schartel and the staff. and Jean Swank. Evan Smith. and theRXTE staff for executing these observations.," We thank Norbert Schartel and the staff, and Jean Swank, Evan Smith, and the staff for executing these observations."
+ We thank the anonymous referee for a helpful report., We thank the anonymous referee for a helpful report.
+ GM thanks the UK PPARC for support., GM thanks the UK PPARC for support.
+ This work has made use of the tools and services available through HEASARC online service. which is operated by GSFC for NASA.," This work has made use of the tools and services available through HEASARC online service, which is operated by GSFC for NASA."
+fluorescent lines.,fluorescent lines.
+ We assumed (1) the primary (accretion-related) power law has a photon index of 1.7 (consistent with the mean photon index of -<0.1 3CRR sources found by Evansetal. 2006)); and (2) it is attenuated by equatorial column densities of 107? ? (consistent with HERGs measured by Evansetal.2006 and Balmaverdeetal.20061) or 10?! ? (to test the Compton-thick case)., We assumed (1) the primary (accretion-related) power law has a photon index of 1.7 (consistent with the mean photon index of $z<0.1$ 3CRR sources found by \citealt{evans06}) ); and (2) it is attenuated by equatorial column densities of $10^{23}$ $^{-2}$ (consistent with HERGs measured by \citealt{evans06} and \citealt{bal06}) ) or $10^{24}$ $^{-2}$ (to test the Compton-thick case).
+ The choice of inclination angle is problematic: Jones&Wehrle(2002) used the jet-counterjet brightness ratio to conclude that the angle to the line of sight. 0. of the jet is 45. but this is different to that subtended by the dusty disk (767) (Ferrarese&Ford 1999)).," The choice of inclination angle is problematic: \cite{jon02} used the jet–counterjet brightness ratio to conclude that the angle to the line of sight, $\theta$, of the jet is $\lappeq$ $^{\circ}$, but this is different to that subtended by the dusty disk $^{\circ}$ ) \citealt{fer99}) )."
+ For illustrative purposes. we adopt an angle of 90° (1.e.. an edge-on geometry) and note that our results are essentially unaffected by the choice of angle until 0<607. after which the reprocessor adopted by MurphyYaqoob(2009) no longer intersects our line of sight.," For illustrative purposes, we adopt an angle of $^{\circ}$ (i.e., an edge-on geometry) and note that our results are essentially unaffected by the choice of angle until $\theta<60^{\circ}$, after which the reprocessor adopted by \cite{my09} no longer intersects our line of sight."
+ The XIS and HXD spectra and models for the 107? em* and 107! 7 cases are shown in Figure 3.., The XIS and HXD spectra and models for the $10^{23}$ $^{-2}$ and $10^{24}$ $^{-2}$ cases are shown in Figure \ref{mytorus_try6and7}.
+ The best-fitting accretion-related power law normalizations were consistent with zero. with 30 upper limits to the 2-10 keV unabsorbed luminosity of 7«10/4 eres ? and |.«10 ergs +. respectively.," The best-fitting accretion-related power law normalizations were consistent with zero, with $\sigma$ upper limits to the 2–10 keV unabsorbed luminosity of $7\times10^{41}$ ergs $^{-1}$ and $4\times10^{42}$ ergs $^{-1}$, respectively."
+ Assuming a black hole mass of 6«105 M.. (Ferrarese&Ford 1999)). we find that any obscured. accretion-related emission ts highly sub-Eddington (Ly/Lraq=10° for Ny=lo? ?; LxíLpga10? for Ny=107! 2).," Assuming a black hole mass of $6\times10^8$ $_{\rm \odot}$ \citealt{fer99}) ), we find that any obscured, accretion-related emission is highly sub-Eddington $L_{\rm X}/L_{\rm Edd}\lappeq10^{-5}$ for $N_{\rm H}$ $10^{23}$ $^{-2}$; $L_{\rm X}/L_{\rm Edd}\lappeq5\times10^{-5}$ for $N_{\rm H}$ $10^{24}$ $^{-2}$ )."
+ The upper limit to the 2-10 keV luminosity in both cases is consistent with the relationship between X-ray luminosity and [μπι luminosity established by Hardeastleetal.(2009)., The upper limit to the 2-10 keV luminosity in both cases is consistent with the relationship between X-ray luminosity and $\mu$ m luminosity established by \cite{hec09}.
+ As an additional test. we considered an extreme Compton thick obseurer CNjj2107?* 2).," As an additional test, we considered an extreme Compton thick obscurer $N_{\rm H}$ $10^{25}$ $^{-2}$ )."
+ Again. this failed to provide a significant improvement to the fit.," Again, this failed to provide a significant improvement to the fit."
+ The upper limit to the 2-10 keV unabsorbed luminosity in this case is 2«10!! eres |., The upper limit to the 2–10 keV unabsorbed luminosity in this case is $2\times10^{44}$ ergs $^{-1}$.
+" However. we note that (1) the peak of the Compton reflection hump lies outside of the usable energy range of the PIN data. meaning that the principal feature used to determine the strength of the accretion-related emission is inaccessible. and (2) a luminosity of 2«10!! eress. ! would be a 6c outlier from the Hardeastleetal.(2009) Lx-Li5,4, relationship."," However, we note that (1) the peak of the Compton reflection hump lies outside of the usable energy range of the PIN data, meaning that the principal feature used to determine the strength of the accretion-related emission is inaccessible, and (2) a luminosity of $2\times10^{44}$ ergs $^{-1}$ would be a $\sigma$ outlier from the \cite{hec09} $L_{\rm X}$ $L_{\rm 15\mu m}$ relationship."
+ Observations with andAstro-H would provide further constraints., Observations with and would provide further constraints.
+ Finally. given the poor signal-to-noise of our PIN data. it is Important to address the known systematic uncertainties of the background. which are of order at the |-sigma level (Fukazawaetal. 2009)).," Finally, given the poor signal-to-noise of our PIN data, it is important to address the known systematic uncertainties of the background, which are of order at the 1-sigma level \citealt{fuk09}) )."
+ In the above analysis. we noticed several points in the PIN band that appear to lie above the model.," In the above analysis, we noticed several points in the PIN band that appear to lie above the model."
+ However. increasing the background by (set using the cornorm parameter in XSPEC) removes any discrepancy. ruling out any evidence of a hard excess given the PIN systematics.," However, increasing the background by (set using the cornorm parameter in ) removes any discrepancy, ruling out any evidence of a hard excess given the PIN systematics."
+ We have presented results from a new 87-ks observation of the prototypical low-excitation radio galaxy 66251., We have presented results from a new 87-ks observation of the prototypical low-excitation radio galaxy 6251.
+ We have shown the following: We wish to thank the anonymous referee for providing constructive comments., We have shown the following: We wish to thank the anonymous referee for providing constructive comments.
+ DAE gratefully acknowledges financial support for this work from NASA under grant number NNXIIAD34G. This research has made use of data obtained from the Suzaku satellite. a collaborative mission between the space agencies of Japan (JAXA) and the USA (NASA).," DAE gratefully acknowledges financial support for this work from NASA under grant number NNX11AD34G. This research has made use of data obtained from the Suzaku satellite, a collaborative mission between the space agencies of Japan (JAXA) and the USA (NASA)."
+ This research has also used the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with NASA.," This research has also used the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with NASA."
+model the bright spot ingress and egress features in this system satisfactorily and results in a poor fit.,model the bright spot ingress and egress features in this system satisfactorily and results in a poor fit.
+" We have thus adapted the model to account for a more complex bright spot by adding four new parameters, following ?,, bringing the total number of variables to 14."," We have thus adapted the model to account for a more complex bright spot by adding four new parameters, following \citet{copperwheat2010}, bringing the total number of variables to 14."
+" These are: 'The data are not good enough to determine the white dwarf limb-darkening coefficient, U,,, accurately."," These are: The data are not good enough to determine the white dwarf limb-darkening coefficient, $U_{w}$, accurately."
+" To find an appropriate limb-darkening coefficient, we follow the procedure outlined in ?,, whereby an estimate of the white dwarf effective temperature and mass is obtained from a first iteration of the fitting process outlined below, assuming a limb-darkening coefficient of 0.345."," To find an appropriate limb-darkening coefficient, we follow the procedure outlined in \citet{littlefair2007}, whereby an estimate of the white dwarf effective temperature and mass is obtained from a first iteration of the fitting process outlined below, assuming a limb-darkening coefficient of 0.345."
+" ? show that typical uncertainties in U,, are ~596,, which leads to uncertainties in Ry/a of ~196.."," \citet{littlefair2007} show that typical uncertainties in $U_{w}$ are $\sim5$ , which leads to uncertainties in $R_{w}/a$ of $\sim1$."
+ These errors have negligible impact on our final system parameters., These errors have negligible impact on our final system parameters.
+" As well as the parameters described above, the model also provides an estimate of the flux contribution from the white dwarf, bright spot, accretion disc and donor."," As well as the parameters described above, the model also provides an estimate of the flux contribution from the white dwarf, bright spot, accretion disc and donor."
+" The white dwarf temperature and distance are found by fitting the white dwarf fluxes from our model to the predictions of white dwarf model atmospheres (?),, as shown in Fig. 3.."," The white dwarf temperature and distance are found by fitting the white dwarf fluxes from our model to the predictions of white dwarf model atmospheres \citep{bergeron1995}, as shown in Fig. \ref{fig:wd_colours}."
+" We find that with the exception of CTCV 2354, all of the systems analysed lie near, or within, the range of white dwarf colours allowed by the atmosphere models of ?,, although the systems do not always lie near the track for the appropriate mass and radius of the white dwarf."," We find that with the exception of CTCV 2354, all of the systems analysed lie near, or within, the range of white dwarf colours allowed by the atmosphere models of \citet{bergeron1995}, although the systems do not always lie near the track for the appropriate mass and radius of the white dwarf."
+ ? compare the temperatures derived using light curve fits to those found using SDSS spectra and GALEX (Galaxy Evolution Explorer) fluxes for a small number of systems and conclude their white dwarf temperatures are accurate to ~1000K. The systems examined by Littlefair et al. (, \citet{littlefair2008} compare the temperatures derived using light curve fits to those found using SDSS spectra and GALEX (Galaxy Evolution Explorer) fluxes for a small number of systems and conclude their white dwarf temperatures are accurate to $\sim$ 1000K. The systems examined by Littlefair et al. (
+2008) are all found to lie close to the Bergeron tracks; it is likely that systems that lie far from the tracks are less accurate.,2008) are all found to lie close to the Bergeron tracks; it is likely that systems that lie far from the tracks are less accurate.
+ We note our temperatures have larger uncertainties than those of ?.., We note our temperatures have larger uncertainties than those of \citet{littlefair2008}.
+ This is because our temperatures take into account the uncertainty in white dwarf mass when comparing the white dwarf fluxes to the models of ?.., This is because our temperatures take into account the uncertainty in white dwarf mass when comparing the white dwarf fluxes to the models of \citet{bergeron1995}.
+" It is possible that our white dwarf colours are affected by contamination from the disc or bright spot, or an unmodelled light source such as a boundary layer."," It is possible that our white dwarf colours are affected by contamination from the disc or bright spot, or an unmodelled light source such as a boundary layer."
+" If our white dwarf colours are incorrect, then our derived white dwarf temperatures will be affected."," If our white dwarf colours are incorrect, then our derived white dwarf temperatures will be affected."
+ Changing the white dwarf temperature will alter ζω., Changing the white dwarf temperature will alter $U_{w}$.
+" Our model fitting measures FR,/a and uses a mass-radius relationship to infer M,,, which is then used to find the mass of the donor star."," Our model fitting measures $R_{w}/a$ and uses a mass-radius relationship to infer $M_{w}$ , which is then used to find the mass of the donor star."
+" However, Uw and Ry are partially degenerate, so Uw therefore affects Ry and My."," However, $U_{w}$ and $R_{w}$ are partially degenerate, so $U_{w}$ therefore affects $R_{w}$ and $M_{w}$."
+ Mw is also affected by temperature changes because the white dwarf mass-radius relationship is temperature dependent., $M_{w}$ is also affected by temperature changes because the white dwarf mass-radius relationship is temperature dependent.
+" The white dwarf temperature also affects the luminosity of the system, and hence distance estimate."," The white dwarf temperature also affects the luminosity of the system, and hence distance estimate."
+ It is therefore important to quantify the effect that incorrect white dwarf temperatures may have on distance estimates and our final derived system parameters., It is therefore important to quantify the effect that incorrect white dwarf temperatures may have on distance estimates and our final derived system parameters.
+" To do this, we altered the white dwarf temperature by 2000K and performed the fitting procedure described above on our best quality, white-dwarf dominated systems."," To do this, we altered the white dwarf temperature by 2000K and performed the fitting procedure described above on our best quality, white-dwarf dominated systems."
+" For lower quality data, the random errors dominate over anysystematic errors, and thus changes to the bestquality data represent a worst case scenario."," For lower quality data, the random errors dominate over anysystematic errors, and thus changes to the bestquality data represent a worst case scenario."
+ We find that changing the white, We find that changing the white
+coupled dark energy as compared to ACDAL. We have therefore shown that both constant. aid variable coupling models. also for cillerent scalar potentials. enhance at any cosmological epoch the cumulative halo mass function | and consequently the probability to detect halos ol any given mass in volume limited surveys with respect to the standard AC'DAL cosmology.,"coupled dark energy as compared to $\Lambda $ We have therefore shown that both constant and variable coupling models, also for different scalar potentials, enhance at any cosmological epoch the cumulative halo mass function – and consequently the probability to detect halos of any given mass in volume limited surveys – with respect to the standard $\Lambda$ CDM cosmology."
+ Future detection of massive clusters at high redshift can therefore be used to disentangle a cosmological constant from dynamical coupled dark energy nmoclels., Future detection of massive clusters at high redshift can therefore be used to disentangle a cosmological constant from dynamical coupled dark energy models.
+also derived.,also derived.
+ Discussion aud conclusion are preseuted iu Section 7., Discussion and conclusion are presented in Section 7.
+ The svuthetic P-L relatious adopted iu this paper are based ou extensive and detailed sets of nonlinear. oilsatiou models including a non-local. time-dependet reatinent of tlk| coupling between pulsation and convection.," The synthetic P-L relations adopted in this paper are based on extensive and detailed sets of nonlinear, pulsation models including a non-local, time-dependent treatment of the coupling between pulsation and convection."
+" These models allow us to predict not ouly he periods and the blue boundary of the instability strip. but also the pulsation amplitudes. the detailed light aud radial velocity curve norphoogy, and the complete opologv of the strip. includiug the red edge (Bonoctal.1999:Fiorentinoetal.2002:Marconi2005. 2010)."," These models allow us to predict not only the periods and the blue boundary of the instability strip, but also the pulsation amplitudes, the detailed light and radial velocity curve morphology, and the complete topology of the strip, including the red edge \citep{bon99,fio02,mar05,marc10}."
+". For each chemica conrposition a1 lnass, an evolutionary nass-Iuuiuositv (ALL) relation was adopted. (seeMaur-conietal.2005.fordetails) aud a wide rauge of effective eniperature was explored."," For each chemical composition and mass, an evolutionary mass-luminosity (M-L) relation was adopted \citep[see][for details]{mar05} and a wide range of effective temperature was explored."
+ We rote that even if the effect of varving he ALE relations has been investigated Or specific chemical compositions (sec.forexamples.Bonoetal.1999.2000:Caputo2005) in this analysis we assume for all the chemical composition a canonical AI-L relation. ucelecting both mass loss aud overshooting during the previous Z7 burning pliase.," We note that even if the effect of varying the M-L relations has been investigated for specific chemical compositions \citep[see, for examples,][]{bon99,bon00,cap05} in this analysis we assume for all the chemical composition a canonical M-L relation, neglecting both mass loss and overshooting during the previous $H$ burning phase."
+ This is a lanitation of he adopted model sets as the above Πο pheuourena affect the P-L relatious aud the correspouding distance determünatious (see Section 7)., This is a limitation of the adopted model sets as the above mentioned phenomena affect the P-L relations and the corresponding distance determinations (see Section 7).
+ Iu total. 17 sets o λος) with varving helm (Y) and netal CZ) abuudaice are considered in this paper. most of them are the saue model sets preseuted in Donoctal. (2010).," In total, 17 sets of model with varying helium $Y$ ) and metal $Z$ ) abundance are considered in this paper, most of them are the same model sets presented in \citet{bon10}."
+. From the resultlue heoretical mstabilitv strips aud he relatious courecting the periods to the intrinsic stellar parameters. svuthetic P-L relations lave becn COSructed.," From the resulting theoretical instability strips and the relations connecting the periods to the intrinsic stellar parameters, synthetic P-L relations have been constructed."
+ To this prrpose. we populated the predicted instability strip by acliptiug the procedure sugeested by Ikeunicuttetal.(1998).," To this purpose, we populated the predicted instability strip by adopting the procedure suggested by \citet{ken98}."
+. Iu. particular. e1000 pulsators were uniforiulv distributed from the blue to the red boundary of the iustabilitv strip .owith a ass law as given by dodi=nP? over the nijass range 5 - 11M. (sceCaputoetal.2000.forfurtherdoetails)..," In particular, $\sim 1000$ pulsators were uniformly distributed from the blue to the red boundary of the instability strip, with a mass law as given by $\mathrm{d}n/\mathrm{d}m = m^{-3}$ over the mass range $5$ - $11 M_{\odot}$ \citep[see][for further details]{cap00}."
+ Iun order to translate the pulsational properties of the investigated Cepheid models iu theSpitzer IRAC bands. we have directly couvolved the predicted 1volometiic elt curves with theSpitzer filter profiles usine the ecucral iuteeral equation (sec.forcxample.Cürardietal.2002):: where 94 is the IRAC Spectral RespouseCurvel. fy is the stellar flux (that correspouds to model atinosplieres of known (Dupe.[M/II]. logq)). FR is the model spectrm of Vega.," In order to translate the pulsational properties of the investigated Cepheid models in the IRAC bands, we have directly convolved the predicted bolometric light curves with the filter profiles using the general integral equation \citep[see, for example,][]{gir2002}: where $S_\lambda$ is the IRAC Spectral Response, $f_\lambda$ is the stellar flux (that corresponds to model atmospheres of known $T_{\mathrm{eff}}, [M/H], \log g $ )), $f_{\lambda}^0$ is the model spectrum of Vega."
+ Concerning the model atmospheres. we have adopted the homogeneous set of updated ATLAS Iurucz model atmospheres and svuthetic fiuxes. ODF ," Concerning the model atmospheres, we have adopted the homogeneous set of updated ATLAS9 Kurucz model atmospheres and synthetic fluxes (new-ODF ."
+models)?.. Figure L shows an example of the svuthetic IRAC han P-L relatious from one of the mocel sets described iu the previous section., Figure \ref{pl_example} shows an example of the synthetic IRAC band P-L relations from one of the model sets described in the previous section.
+ All the svutietic P-L relations were fitted with the form of Appa:€(|bxiP). restrictec for pulsators within the period range ofloge 0.1€logtP?)<2.0 (asinBonoetal.2010).," All the synthetic P-L relations were fitted with the form of $M_{\mathrm{IRAC}}=a+b\times \log(P)$, restricted for pulsators within the period range of $0.4 \leq \log(P) \leq 2.0$ \citep[as in][]{bon10}."
+". The fitted P-L slopes (5) iux intercepts (a) for cach of the iioels were ΠΠ iu Table P. and 2.. respectively,"," The fitted P-L slopes $b$ ) and intercepts $a$ ) for each of the models were summarized in Table \ref{plslope} and \ref{plzp}, respectively."
+ A majority of the mode sets do not have pulsators with ος}«0.L, A majority of the model sets do not have pulsators with $\log(P)<0.4$.
+ For a cw of the model sets which have a sia] uuniber of short-)iod pulsators. the differences of he P-L slopes zx intercepts between the P-L relations derived frou ful )xiod range aud those eiven in Tab eslo and 2 do not exceed 0.01.," For a few of the model sets which have a small number of short-period pulsators, the differences of the P-L slopes and intercepts between the P-L relations derived from full period range and those given in Tables \ref{plslope} and \ref{plzp} do not exceed $0.01$."
+ In Figure 2.. the slopes «f the svuthetic P-L relations in various bands were couxuwed for six of the oulsating model sets (or chemical compositions). where he linear GVEJN P-L slopes were adopted from Table 2 of Bonoetal.(2010)(2010).," In Figure \ref{pl_multiband}, the slopes of the synthetic P-L relations in various bands were compared for six of the pulsating model sets (or chemical compositions), where the linear $BVIJK$ P-L slopes were adopted from Table 2 of \citet{bon10}."
+.. As expected. the slopes of the P-L relation monotouically decrease from B. to A baud (see.forexample.Madore&Freediuan1991: 2008).. and “fatten out” in the mid-infrared.," As expected, the slopes of the P-L relation monotonically decrease from $B$ to $K$ band \citep[see, for example,][]{mad91,ber96,cap00,fio02,fio07,fre08,nge08}, and “flatten out” in the mid-infrared."
+ However. he L5gan P-L relations show a shel: lucrease in thei slopes when compared to the “fatter” slopes defined roni 3.6401 aud 8.0724 baud P-L slopes.," However, the $4.5\mu{\mathrm m}$ P-L relations show a slight increase in their slopes when compared to the “flatter” slopes defined from $3.6\mu{\mathrm m}$ and $8.0\mu{\mathrm m}$ band P-L slopes."
+" This slight increase of the [5,140 P-L slopes. and i some extent o the Sayan baud P-L slopes. may be explained due o the presence of CO absorption features shown in the tian to Gyan spectral region"," This slight increase of the $4.5\mu{\mathrm m}$ P-L slopes, and in some extent to the $5.8\mu{\mathrm m}$ band P-L slopes, may be explained due to the presence of CO absorption features shown in the $\sim4\mu{\mathrm m}$ to $\sim6\mu{\mathrm m}$ spectral region"
+" ).. 3. τοῦ, 120.. 564.",radio loud sources.
+. 4051.. 1314... 279.. 3783... G24.. and 5548)) included in both our aad samples.," There are 22 objects, and ) included in both our and samples."
+ Figure 3a plots their N-vav Iuminosities againstNALA ones., Figure 3a plots their X-ray luminosities against ones.
+ Clearly. (here is no svslemalie bias between the observed. luminosities by different instruments. however. ὃ out of 22 sources show variation wilh amplitude above 1.5.," Clearly, there is no systematic bias between the observed luminosities by different instruments, however, 8 out of 22 sources show variation with amplitude above 1.5."
+ The relationship of line EW observed by audNALAL is plotted in Figure 3b., The relationship of line EW observed by and is plotted in Figure 3b.
+ We can see that most of EW observed by ave consistent with those observed byNALA., We can see that most of EW observed by are consistent with those observed by.
+ Considering that data has better spectral resolution in the Fe Ka band. which is essential to measure (he narrow line. we directly. adopt data for these 22 sources aud dropXA ones in the combined sample.," Considering that data has better spectral resolution in the Fe $\alpha$ band, which is essential to measure the narrow line, we directly adopt data for these 22 sources and drop ones in the combined sample."
+ Fieure 4 plots the correlation between the EW ancl luminosity for all objects in the large sample (a=—0.2015+0.0426. Rs = —0.469) and Lor radio-quiet objects oulv fa=—0.1019€0.0524. Rs = —0.266).," Figure 4 plots the correlation between the EW and luminosity for all objects in the large sample $\alpha=-0.2015\pm0.0426$, Rs = $-0.469$ ) and for radio-quiet objects only $\alpha=-0.1019\pm0.0524$, Rs = $-0.266$ )."
+ The fitting results are consistent with those derived from andNALA data respectively., The fitting results are consistent with those derived from and data respectively.
+ The values for three dillerent samples are listed together in Table 3.., The values for three different samples are listed together in Table \ref{tbl-3}.
+ Note for the Chandra only sample. (he difference between (he correlation index from RQ+KRL sources and RQ sources only is not obvious as for NMM and combined samples.," Note for the Chandra only sample, the difference between the correlation index from RQ+RL sources and RQ sources only is not obvious as for XMM and combined samples."
+ This is mainly because that there is very limited number of luminous sources in the Chandra only sample: only one RQ sources (IL 18214643) shows Ly above LOY? erg Jl thus the correlation index is extremely sensitive to its EW measurement: for I1 13212-6423 we fixed the line width at 1000 km ! to get an EW of 40 eV: while the best fit value gives a line width of 10000 km ! and an EW of 140 eV: adopting the later values we get a linear correlation index of —0.1062+0.0513.," This is mainly because that there is very limited number of luminous sources in the Chandra only sample: only one RQ sources (H 1821+643) shows $_X$ above $^{45.5}$ erg $^{-1}$, thus the correlation index is extremely sensitive to its EW measurement: for H 1821+643 we fixed the line width at 1000 km $^{-1}$ to get an EW of 40 eV; while the best fit value gives a line width of 10000 km $^{-1}$ and an EW of 140 eV; adopting the later values we get a linear correlation index of $-0.1062\pm0.0573$."
+ We can clearly see (hat when the RL sources were excluded [rom the sample the correlation between the line line EW and X-ray huminosity became much weaker (with a confidence level less (han 20)., We can clearly see that when the RL sources were excluded from the sample the correlation between the line line EW and X-ray luminosity became much weaker (with a confidence level less than $\sigma$ ).
+ To show the results more clearly. in Figure 5 we plot the mean EW for 5 luminosity bins for all AGNs and just the radio-quiet objects in the combined sample.," To show the results more clearly, in Figure 5 we plot the mean EW for 5 luminosity bins for all AGNs and just the radio-quiet objects in the combined sample."
+ Following Page et al. (, Following Page et al. (
+2004a). the upper limits were taken to be half of (he value. together with an equally sized error bar.,"2004a), the upper limits were taken to be half of the value, together with an equally sized error bar."
+ Similar (o Page et al..," Similar to Page et al.,"
+ we find a clear anti-correlation for RQ-RL sources (Figure 5a)., we find a clear anti-correlation for RQ+RL sources (Figure 5a).
+ However. for RQ only sources. we see no anti-correlation in Figure 5b.," However, for RQ only sources, we see no anti-correlation in Figure 5b."
+ suggestions for the origin of the neutral Fe Na line include the putative molecular torus and/or the DLR. (broad-line region)., Suggestions for the origin of the neutral Fe $\alpha$ line include the putative molecular torus and/or the BLR (broad-line region).
+ Since produced al a much larger scale. (he line should," Since produced at a much larger scale, the line should"
+We have shown that the WD IFMR places powerful constraints on the energy released during the TP-AGB stage of stellar evolution.,We have shown that the WD IFMR places powerful constraints on the energy released during the TP-AGB stage of stellar evolution.
+ Future observations of TP-AGB populations will constrain the helium yields of these stars- an important factor for chemical evolution models to consider., Future observations of TP-AGB populations will constrain the helium yields of these stars- an important factor for chemical evolution models to consider.
+" Alternatively, nominal assumptions regarding the yield of these stars, in conjunction with this work, result in a narrower range of possible TP-AGB population luminosity than previously considered."," Alternatively, nominal assumptions regarding the yield of these stars, in conjunction with this work, result in a narrower range of possible TP-AGB population luminosity than previously considered."
+ We would like to thank David Weinberg for his encouragement and thought-provoking discussion., We would like to thank David Weinberg for his encouragement and thought-provoking discussion.
+not observed to correlate with the rotational parameter.,not observed to correlate with the rotational parameter.
+ 'This lack of correlation was already found by Benderetal. (1994)., This lack of correlation was already found by \cite{bender94}.
+". However, the distribution present a wide spread of positive ha values at any Όγοι/σ. indicating a significant rotationally-supported disc component."," However, the distribution present a wide spread of positive $h_{4}$ values at any $v_{rot}/\sigma$, indicating a significant rotationally-supported disc component."
+" We find that almost 85 per cent of our sample galaxies are characterised by a cold stellar component extending out to the outermost regions reached by our study, i.e. 3re."," We find that almost 85 per cent of our sample galaxies are characterised by a cold stellar component extending out to the outermost regions reached by our study, i.e. $\sim 3 r_{e}$."
+" The stellar velocity dispersion is observed to quickly decline with radius, such that ordered motion prevails over random stellar orbits."," The stellar velocity dispersion is observed to quickly decline with radius, such that ordered motion prevails over random stellar orbits."
+ The galaxies are well described by oblate rotationally supported models., The galaxies are well described by oblate rotationally supported models.
+" The results of our analysis confirms previous circumstantial findings by Bender&Ni-eto (1990),, Daviesetal. (1983),, and Hallidayetal.(2001)."," The results of our analysis confirms previous circumstantial findings by \cite{bender90}, \cite{davies83}, and \cite{halliday01}."
+". However, unlike these earlier studies, we are able to spatially resolved enough apertures at large galactocentric radii (i.e., r2 0.5re), where the data have a sufficiently high S/N to make robust measurements."," However, unlike these earlier studies, we are able to spatially resolved enough apertures at large galactocentric radii (i.e., $r \geq 0.5 r_{e}$ ), where the data have a sufficiently high S/N to make robust measurements."
+" During an early star-forming collapse (e.g., Chiosi&Carraro 2002;; Kawata&Gibson 2003)), a cold and extended stellar component can naturally originate from the high angular momentum gas situated in the outer regions of the forming galaxies (e.g., Bekki1998;; Rixetal. 1999;; Naab&Burkert 2001;; Kawata&Gibson 2003))."," During an early star-forming collapse (e.g., \citealt{chiosi02}; \citealt{kawata03}) ), a cold and extended stellar component can naturally originate from the high angular momentum gas situated in the outer regions of the forming galaxies (e.g., \citealt{bekki98}; \citealt{rix99}; \citealt{naab01}; \citealt{kawata03}) )."
+" On the other hand, the gas located at inner radii tend to sink toward the galactic centre where it forms new stars."," On the other hand, the gas located at inner radii tend to sink toward the galactic centre where it forms new stars."
+" This outside-in formation scenario (e.g., Pipinoetal. 2006b)) is predicted to induce a shallow positive age gradient in the galactic stellar population."," This outside-in formation scenario (e.g., \citealt{pipino06b}) ) is predicted to induce a shallow positive age gradient in the galactic stellar population."
+ In Paper II we find that the 14 low-luminosity galaxies have an average age gradient of 0.06+0.13 dex per decade., In Paper II we find that the 14 low-luminosity galaxies have an average age gradient of $0.06 \pm 0.13$ dex per decade.
+" The gradients are such that central regions (i.e., r<0.5re) are a few Gyr younger with respect to the outer parts."," The gradients are such that central regions (i.e., $r \leq 0.5r_{e}$ ) are a few Gyr younger with respect to the outer parts."
+" The stellar age of these outer galactic regions is on average older than 8 Gyr (see table 3 in Paper II) In other words, the cold stellar component of the sample galaxies is old and it was formed at high redshift, i.e. z21."," The stellar age of these outer galactic regions is on average older than 8 Gyr (see table 3 in Paper II) In other words, the cold stellar component of the sample galaxies is old and it was formed at high redshift, i.e. $z \geq1$."
+" 'The radial variation in the ellipticity of the isophotes, and their prevalent discyness, observed in our data sample could be attributed to the central concentration of stars"," The radial variation in the ellipticity of the isophotes, and their prevalent discyness, observed in our data sample could be attributed to the central concentration of stars"
+ ↕↕∨⊞∆↕∖↧⊳∖∖⋎↓↥↕≼⇍↓↕⊲↓⊔∩⊾⋏∙≟↓⋅⋜⋯⋅⊳∖↿↓⊔⋅⊔↓⋜↧↓≻∖⇁⋜↧↓⋯⊾⊲↓⊔⋜↧⊳∖↓≻∢⊾≼⋰↓∐≼∼↓⋅∢⊾≼∼↿⋜⋯⋏∙≟↓∢⊾⊳ while for all the other parameters (major axis. minor axis and position angle) we used the values obtained: with the Gaussian fit.,", which integrates the map value in a specific rectangle, while for all the other parameters (major axis, minor axis and position angle) we used the values obtained with the Gaussian fit."
+ For irregular resolved. sources the total [Lux density was calculated: using the taskTVSTAT which allows us to use an irregular area to integrate the map value., For irregular resolved sources the total flux density was calculated using the task which allows us to use an irregular area to integrate the map value.
+ Detected sources separated by less than two times the value of our svnthesized beam size (e. 25 aresec) and with a us ratio lower than 2 have been considered. as an unique source., Detected sources separated by less than two times the value of our synthesized beam size $i.e.$ 25 arcsec) and with a flux ratio lower than 2 have been considered as an unique source.
+ We adopted this criterion because the component Hux density ratio of physically doubles is usually small 2) while the projection pairs can have arbitrarily large. Dux density ratio (Condon. Condon and Hazard. 1982).," We adopted this criterion because the component flux density ratio of physically doubles is usually small $\lsimeq$ 2) while the projection pairs can have arbitrarily large flux density ratio (Condon, Condon and Hazard, 1982)."
+ Considering all the available observations we detected a total of SGT sources at ο 5 m level (44 of which have multiple components) over a total area of 4.222 deg?., Considering all the available observations we detected a total of 867 sources at $\geq$ 5 $\sigma$ level (44 of which have multiple components) over a total area of 4.222 $^2$.
+" The catalogue with all the S67 sources (921 components) reports the name of the source. the peak flux density Sp in mv. the total lux density 5; in mJy. the RA and DEC (J2000). the full width half maximum (ENIIN) of the major and minor axes 03, and 6,, (in aresec). the positional angle PA of the major axis (in degrees) and the olf-axis values in the VLA map (in arcmin)."," The catalogue with all the 867 sources (921 components) reports the name of the source, the peak flux density $_P$ in mJy, the total flux density $_I$ in mJy, the RA and DEC (J2000), the full width half maximum (FWHM) of the major and minor axes $\theta_M$ and $\theta_m$ (in arcsec), the positional angle PA of the major axis (in degrees) and the off-axis values in the VLA map (in arcmin)."
+ The dillerent components of multiple sources are labeled “A. 7. ete.," The different components of multiple sources are labeled “A”, “B”, etc.,"
+ followed by a line labeled. 7I in which Εαν ancl position for the total sources are given., followed by a line labeled “T” in which flux and position for the total sources are given.
+ For these total sources the position have been computed as the —ux-weighted average position for all the components., For these total sources the position have been computed as the flux-weighted average position for all the components.
+ Table 3 shows the first page of the catalogue as an example., Table 3 shows the first page of the catalogue as an example.
+ In Table 40 we report the number of radio sources etected in NI. N2 and N3. while in Figure 7. we show the istribution of the peak Mux density aid the total to peak —ux ratio as a function of peak Ilux for all the S67 sources.," In Table \ref{sources_tab}
+ we report the number of radio sources detected in N1, N2 and N3, while in Figure \ref{flux_ratio} we show the distribution of the peak flux density and the total to peak flux ratio as a function of peak flux for all the 867 sources."
+ ο.'ontour maps of allthe 44 double or multiple sources are gajown in Figure S 9 and 10., Contour maps of allthe 44 double or multiple sources are shown in Figure \ref{double_contour1} \ref{dN2} and \ref{dN3}. .
+the instantaneous luminosity at z—1.0 (red open circles).,the instantaneous luminosity at $z=1.0$ (red open circles).
+ The best-fit slope for the instantaneous luminosity is (0.95+0.12)., The best-fit slope for the instantaneous luminosity is $(0.95 \pm 0.12)$.
+ As we go to higher redshifts the AGN have ]uminosities closer to their peak luminosities and hence we see a higher mean value and a tighter correlation 55)., As we go to higher redshifts the AGN have luminosities closer to their peak luminosities and hence we see a higher mean value and a tighter correlation 5).
+ The difference between the peak and the instantaneous luminosity is most prominent at higher mass halos., The difference between the peak and the instantaneous luminosity is most prominent at higher mass halos.
+ We associate this suppression with feedback from AGN., We associate this suppression with feedback from AGN.
+" In Paper I, we used the Man—c relation to show that black holes residing in higher mass halos enter the feedback dominated phase at an earlier time than black holes populating lower mass halos."," In Paper I, we used the $M_{\rm{BH}} -\sigma$ relation to show that black holes residing in higher mass halos enter the feedback dominated phase at an earlier time than black holes populating lower mass halos."
+ T'hus feedback effects will alter the Mnaio—Lpoi correlation resulting in a wider distribution of host halo masses for a given AGN luminosity at lower redshift., Thus feedback effects will alter the $M_{\rm{halo}}-L_{\rm{Bol}}$ correlation resulting in a wider distribution of host halo masses for a given AGN luminosity at lower redshift.
+ This result is also in agreement with ? who conclude that the peak luminosity is more correlated with halo mass than the instantaneous luminosity of AGN and hence a better indicator of clustering., This result is also in agreement with \citet{lidzetal06} who conclude that the peak luminosity is more correlated with halo mass than the instantaneous luminosity of AGN and hence a better indicator of clustering.
+" However, we cannot measure the peak luminosity of the AGN in practice."," However, we cannot measure the peak luminosity of the AGN in practice."
+ We model the mean occupation function of AGN in dark, We model the mean occupation function of AGN in dark
+Iu this paper. we male the first perturbative prediction for the A power spectrum.,"In this paper, we make the first perturbative prediction for the $A$ power spectrum."
+ There are reasons to think that perturbation theory for A might work even better than for à., There are reasons to think that perturbation theory for $A$ might work even better than for $\delta$.
+ At small scales. the variance in lis much smaller than in 6 in the nonlinear regine.," At small scales, the variance in $A$ ismuch smaller than in $\delta$ in the nonlinear regime."
+ Also. it is tempting to try to exploit the simplicity of ((1)) iu a Lagrangian approach.," Also, it is tempting to try to exploit the simplicity of \ref{eqn:logcont}) ) in a Lagrangian approach."
+ However. working with A also presents some challenges.," However, working with $A$ also presents some challenges."
+" For the ó field. smoothing (outo a erid. for example) has fairly trivial effects: at large scales. it leaves phases intact aud dampens amplitudes neeligibly,"," For the $\delta$ field, smoothing (onto a grid, for example) has fairly trivial effects; at large scales, it leaves phases intact and dampens amplitudes negligibly."
+ But smoothing is crucial to consider for 4. giviug different large-scale biases for different. noothing scales (Nevrineletal.2009).," But smoothing is crucial to consider for $A$ , giving different large-scale biases for different smoothing scales \citep{NSS09}."
+. It is quite difficult to model finalcoucitious Eulerian smoothing using Lagrangian variables: this is why we use an Eulerian approach., It is quite difficult to model final-conditions Eulerian smoothing using Lagrangian variables; this is why we use an Eulerian approach.
+ For this. we use a Tavlor series expansion of the logarithin: the slow couverecuce of this alternating series iu itself porteuds low accuracy at highly nonlinear scales where 9 is larec.," For this, we use a Taylor series expansion of the logarithm; the slow convergence of this alternating series in itself portends low accuracy at highly nonlinear scales where $\delta$ is large."
+ Despite these challenges. the perturbative calculations we preseut below achieve good agreement with simulation measurements.," Despite these challenges, the perturbative calculations we present below achieve good agreement with simulation measurements."
+ Iu the next section. we review the basic framework of Enlerian dvuamics aud cliscuss the logaritlanic transformation.," In the next section, we review the basic framework of Eulerian dynamics and discuss the logarithmic transformation."
+ We calculate the power spectrum of the A field in standard perturbation theory in section 3. and iu renormalized perturbation theory im section [.," We calculate the power spectrum of the $A$ field in standard perturbation theory in section 3, and in renormalized perturbation theory in section 4."
+ Iu section 5. we prescut our calculation aud compare it with simulation. and we conclude in section 6.," In section 5, we present our calculation and compare it with simulation, and we conclude in section 6."
+ The gravitational dvunamics of a pressureless fluid before shell crossing is governed by the contiuuitv. Euler. aud Poisson equations.," The gravitational dynamics of a pressureless fluid before shell crossing is governed by the continuity, Euler, and Poisson equations."
+" Tere. H= dluetr)ídluristhe IIubble expansion rate. e(7) is the scale factor. O,,(7) is the ratio of matter density to critical density. and Py(x.7) is the Newtonian potential."," Here, $\mH= d\ln a(\tau)/ d\ln\tau$ is the Hubble expansion rate, $a(\tau)$ is the scale factor, $\Omega_m(\tau)$ is the ratio of matter density to critical density, and $\Phi_N(\vx,\tau)$ is the Newtonian potential."
+ In principle. perturbation theory of the loe-deusity field is possible using the loe-coutinuity equation. ((1)). together with the last two equations in ((2)).," In principle, perturbation theory of the log-density field is possible using the log-continuity equation, \ref{eqn:logcont}) ), together with the last two equations in \ref{eqn:dyneqn_sep}) )."
+ However. we found it more straightforward to take smoothing into account if we expand 4 in terms of à. so this is the approach that we tale.," However, we found it more straightforward to take smoothing into account if we expand $A$ in terms of $\delta$, so this is the approach that we take."
+ First we briefly review results for the conventional overdensity field 9., First we briefly review results for the conventional overdensity field $\delta$.
+" Following Crocce&Scoccimarre (20062). the equation of motion in Fourier space can be written in à compact form by defining the two-compoucut variable The index aC[p.c] stands for density or velocity variables. respectively,"," Following \cite{CS06a,CS06b}, , the equation of motion in Fourier space can be written in a compact form by defining the two-component variable The index $a\in\{\rho, v\}$ stands for density or velocity variables, respectively."
+" Where umnberced iudices are more convenicut. we also use &C(1.2). Wy=V, and VY,=V,."," Where numbered indices are more convenient, we also use $a\in\{1, 2\}$, $\Psi_1=\Psi_{\rho}$, and $\Psi_2=\Psi_v$."
+ The equation of motion then reads with the convention that repeated Fourier arguments are integrated over., The equation of motion then reads with the convention that repeated Fourier arguments are integrated over.
+ The time 4j=lie(r) iu a DEiustein«de Sitter (EdS) universe., The time $\eta=\ln a(\tau)$ in a Einstein-de Sitter (EdS) universe.
+ Hore the coustaut matrix derived for an EdS universe. is still applicable iu other cosimnologies with negligible corrections to the coefficients.using y=luD(z) (ith D the growth factor).," Here the constant matrix derived for an EdS universe, is still applicable in other cosmologies with negligible corrections to the coefficients,using $\eta=\ln
+D(\tau)$ (with $D$ the growth factor)."
+ The svuunetrized vertex matrix 5454.19 giveu by uio(k.ki.ky)=syo(k.κ». ki). aud 3—0 otherwise.," The symmetrized vertex matrix $\gamma_{abc}$ is given by $\gamma_{112}(\vk,\vk_1,\vk_2)=\gamma_{121}(\vk,\vk_2,\vk_1)$ , and $\gamma=0$ otherwise."
+" Then the formal integral solution to ((1)) cau bederived as where o,(sx) denotes the initial condition o,(k)=W,(k.;jj 0). aud the linear propagator ga4í0]) is givenby"," Then the formal integral solution to \ref{eqn:dyneqn}) ) can bederived as where $\phi_a(\vk)$ denotes the initial condition $\phi_a(\vk)\equiv \Psi_a(\vk,\eta=0)$ and the linear propagator $g_{ab}(\eta)$ is givenby"
+Nuclear burning in the degenerate matter of a dense C+O white dwarf. once ignited. is believed to propagate on microscopic scales as a conductive fame. wrinkled and stretched by local turbulence. but with essentially the laminar velocity.,"Nuclear burning in the degenerate matter of a dense C+O white dwarf, once ignited, is believed to propagate on microscopic scales as a conductive flame, wrinkled and stretched by local turbulence, but with essentially the laminar velocity."
+" At high densities near the center of a Chandrasekhar-mass white dwarf the typical length scales for the width of the flame are a fraction of a millimeter,", At high densities near the center of a Chandrasekhar-mass white dwarf the typical length scales for the width of the flame are a fraction of a millimeter.
+ However. this is not the most relevant length scale.," However, this is not the most relevant length scale."
+ Due to the very high Reynolds numbers. which are of the order of {011 on typical macroscopic scales of I0*cm. macroscopic flows are highly turbulent and interact with the flame. in principle down to the Kolmogorov scale of 10.em.," Due to the very high Reynolds numbers, which are of the order of $^{14}$ on typical macroscopic scales of $^7$ cm, macroscopic flows are highly turbulent and interact with the flame, in principle down to the Kolmogorov scale of $^{-3}$ cm."
+ This means that all kinds of hydrodynamic instabilities feed energy into a turbulent cascade. including the buoyancy-driven Rayleigh-Taylor instability and the shear-driven Kelvin-Helmholtz instability (see. e.g. ? and ?)).," This means that all kinds of hydrodynamic instabilities feed energy into a turbulent cascade, including the buoyancy-driven Rayleigh-Taylor instability and the shear-driven Kelvin-Helmholtz instability (see, e.g., \citealt{niemeyer-hillebrandt-95a} and \citealt{hillebrandt-niemeyer-96}) )."
+" Consequently. the picture that emerges is more that of a ""flame brush” spread over the entire turbulent regime rather than a wrinkled flame surface."," Consequently, the picture that emerges is more that of a “flame brush” spread over the entire turbulent regime rather than a wrinkled flame surface."
+ For such a flame brush. the relevant minimum length scale is the so-called Gibson scale. defined as the lower bound for the curvature radius of flame wrinkles caused by turbulent stress.," For such a flame brush, the relevant minimum length scale is the so-called Gibson scale, defined as the lower bound for the curvature radius of flame wrinkles caused by turbulent stress."
+ Thus. if the thermal diffusion scale is much smaller than the Gibson scale (which is the case for the physical conditions of interest here) small segments of the flame surface are unaffected by large scale turbulence and behave as unperturbed laminar flames (“fHlamelets”).," Thus, if the thermal diffusion scale is much smaller than the Gibson scale (which is the case for the physical conditions of interest here) small segments of the flame surface are unaffected by large scale turbulence and behave as unperturbed laminar flames (“flamelets”)."
+ On the other hand side. since the Gibson seale is. at high densities. several orders of magnitude smaller than the integral scale set by the Rayleigh-Taylor eddies and many orders of magnitude larger than the thermal diffusion scale. both transport and burning times are determmed by the eddy turnover times. and the effective velocity of the burning front is independent of the laminar burning velocity.," On the other hand side, since the Gibson scale is, at high densities, several orders of magnitude smaller than the integral scale set by the Rayleigh-Taylor eddies and many orders of magnitude larger than the thermal diffusion scale, both transport and burning times are determined by the eddy turnover times, and the effective velocity of the burning front is independent of the laminar burning velocity."
+ This general picture has indeed been verified by laboratory combustion experiments (see. e.g.. 2)).," This general picture has indeed been verified by laboratory combustion experiments (see, e.g., \citealt{abdelgayed-etal-87}) )."
+ ? and ? have presented a umerical realization of this general concept., \cite{niemeyer-95} and \cite{niemeyer-hillebrandt-95a} have presented a numerical realization of this general concept.
+ Their basic assumption was that wherever one finds turbulence this turbulence is fully developed and homogeneous. ie. the turbulent velocity fluctuations on a length scale / are given by the Kolmogorov law where L is the integral scale. assumed to be equal to the Rayleigh-Taylor scale.," Their basic assumption was that wherever one finds turbulence this turbulence is fully developed and homogeneous, i.e. the turbulent velocity fluctuations on a length scale $l$ are given by the Kolmogorov law where $L$ is the integral scale, assumed to be equal to the Rayleigh-Taylor scale."
+ Following the ideas outlined above. one can also assume that the thickness of the turbulent flame brush on the scale / is of the order of/ itself.," Following the ideas outlined above, one can also assume that the thickness of the turbulent flame brush on the scale $l$ is of the order of $l$ itself."
+With these two assumptions and the definition of the Gibson scale one finds for HatsSISE~Ary and where (έως)Man defines. fips. ων 08 the laminar burning speed and «;(/) is the turbulent flame velocity on the scale 7.,"With these two assumptions and the definition of the Gibson scale one finds for $l_{\text{gibs}} \sil l \sil L \simeq \lambda _{\text{RT}}$ and where $v(l_{\text{gibs}}) = u_{\text{lam}}$ defines $l_{\text{gibs}}$, $u_{\text{lam}}$ is the laminar burning speed and $u_t(l)$ is the turbulent flame velocity on the scale $l$ ."
+ In a second step this model of turbulent combustion ts coupled to a finite volume hydro scheme such as the PPM-code PROMETHEUS (?).., In a second step this model of turbulent combustion is coupled to a finite volume hydro scheme such as the PPM-code PROMETHEUS \citep{fryxell-etal-89}.
+" Since in every finite volume scheme scales smaller than the grid size cannot be resolved. we express ls, in. terms of the grid size A. the (unresolved) turbulent kinetic energy q. and the laminar burning velocity: Here 4 is determined from a sub-grid model (??) and. finally. the effective turbulent velocity of the flame brush on scale A is given by with ¢(A)=2g and tyrX\/yA. where y is the local gravitational acceleration."," Since in every finite volume scheme scales smaller than the grid size cannot be resolved, we express $l_{\text{gibs}}$ in terms of the grid size $\Delta$, the (unresolved) turbulent kinetic energy $q$, and the laminar burning velocity: Here $q$ is determined from a sub-grid model \citep{clement-93,
+niemeyer-hillebrandt-95a} and, finally, the effective turbulent velocity of the flame brush on scale $\Delta$ is given by with $v(\Delta) = \sqrt {2q}$ and $v_{RT} \propto \sqrt {g \Delta}$, where $g$ is the local gravitational acceleration."
+ The numerical scheme outlined so far has been applied to two-dimensional simulations of explosive C-burning in Mey- mass C+O white dwarfs by ? and ?. for a variety of initial conditions.," The numerical scheme outlined so far has been applied to two-dimensional simulations of explosive C-burning in $_{\text{Ch}}$ -mass C+O white dwarfs by \cite{niemeyer-hillebrandt-95a} and \cite{niemeyer-etal-96}
+ for a variety of initial conditions."
+ They found that. with the exception of a model that was ignited off-center in a single small blob. all models were unbound at the end of the computations. and between 0.59 M. (for central ignition) and 0.65 .. of matter had been burnt. which would guarantee an observable. although weak. Type la supernova. but certainly not an average one. in particular because the expansion velocities were too low.," They found that, with the exception of a model that was ignited off-center in a single small blob, all models were unbound at the end of the computations, and between 0.59 $_{\odot}$ (for central ignition) and 0.65 $_{\odot}$ of matter had been burnt, which would guarantee an observable, although weak, Type Ia supernova, but certainly not an average one, in particular because the expansion velocities were too low."
+ There are two possible interpretations of these results., There are two possible interpretations of these results.
+" Either they show that the ""normal"" Type Ia supernova isnor the result of the explosion of a Mcy-mass C+O white dwarf. caused by a deflagration wave. or the modeling of thermonuclear combustion as outlined so far is still insufficient."," Either they show that the “normal” Type Ia supernova is the result of the explosion of a $_{\text{Ch}}$ -mass C+O white dwarf, caused by a deflagration wave, or the modeling of thermonuclear combustion as outlined so far is still insufficient."
+ In this paper we investigate the second possibility and. in particular. present results obtained by means of an improved numerical scheme.," In this paper we investigate the second possibility and, in particular, present results obtained by means of an improved numerical scheme."
+ The main aim is to cure a problem that is common to all finite volume schemes. namely that. although discontinuities are captured automatically. they are smeared artificially over several grid zones.," The main aim is to cure a problem that is common to all finite volume schemes, namely that, although discontinuities are captured automatically, they are smeared artificially over several grid zones."
+ In many applications this broadening of discontinuities might not be a big problem but for modeling turbulent combustion one cannot be satisfied., In many applications this broadening of discontinuities might not be a big problem but for modeling turbulent combustion one cannot be satisfied.
+ The reasons for this are simple., The reasons for this are simple.
+ First. as we discussed already. it is important to reproduce the geometry of the flame front as accurately as possible. because its effective surface area also determines the rate of fuel consumption.," First, as we discussed already, it is important to reproduce the geometry of the flame front as accurately as possible, because its effective surface area also determines the rate of fuel consumption."
+ If we think of the flame in terms of a surface of a given fractal dimension the minimum scale determines. essentially its effective velocity., If we think of the flame in terms of a surface of a given fractal dimension the minimum scale determines essentially its effective velocity.
+ Therefore it is important to resolve the small scales as accurately as possible., Therefore it is important to resolve the small scales as accurately as possible.
+ Secondly. since nuclear reactions are very sensitive to temperature and their rates depend on a very high power of T. mixing fuel and ashes numerically in certain mesh points on both sides of the front and smearing out the temperature gradient there willlead to huge errors in the (local) energy generation rates.," Secondly, since nuclear reactions are very sensitive to temperature and their rates depend on a very high power of $T$ , mixing fuel and ashes numerically in certain mesh points on both sides of the front and smearing out the temperature gradient there willlead to huge errors in the (local) energy generation rates."
+ Therefore.," Therefore,"
+we obtain the following set of Fredholm iutegral equations where ancl kernel functions independent of kiud of primary Exaiples of distribution functions 0Mμυ for EAS electron iud truucated mmo size spectra are preseuted Fie.,"we obtain the following set of Fredholm integral equations where and kernel functions independent of kind of primary Examples of distribution functions $\partial
+W/\partial x$ for EAS electron and truncated muon size spectra are presented Fig."
+ 1.2 respectively.," 1,2 respectively."
+" These data were obtained by CORSIKA code for primary proton (4=1. eiipty sviubols)aud iron (4=56. Ἡled svaubols) nuclei at euerev E=3.2:109 GeV. and zenith augle 0«I8""."," These data were obtained by CORSIKA code for primary proton $A\equiv1$, empty symbols)and iron $A\equiv56$, filled symbols) nuclei at energy $E=3.2\cdot10^6$ GeV and zenith angle $\theta<18^0$."
+ The lines in Fie., The lines in Fig.
+" 1.2 corresspoud to log-Caussian distributious which were counted based oun known values of Weal «ee and variances. σι>AE.0), ]t is seen. hat the distribution functions slightly depend on a kind of a primary nucleus. especially the xight-hanud sides of distributions."," 1,2 correspond to log-Gaussian distributions which were counted based on known values of mean $<x>=1$ and variances $\sigma^2_{e,\mu}(E,\theta)$ It is seen, that the distribution functions slightly depend on a kind of a primary nucleus, especially the right-hand sides of distributions."
+ It is an müportaut fact. because the contribution of the left-hane sides of distributions is negligible small at steep primary cacrey spectra GS(OEE5jm ," It is an important fact, because the contribution of the left-hand sides of distributions is negligible small at steep primary energy spectra $\partial\Im_{A}/\partial E\sim E^{-3}$ "
+"ddata. we cannot draw a firm conclusion on this We also modeled the multi-band spectrum in the ""low"" (2003) and ""high"" (2005) state with a resonant cyclotron scattering model plus a power-law. and we find a good fit in the ""low"" state with 8=0.49(1). KT.=0.36(3). 7=1.2(D) and ΓΞ1.0(1) (see Reaet al.","data, we cannot draw a firm conclusion on this We also modeled the multi-band spectrum in the “low” (2003) and “high” (2005) state with a resonant cyclotron scattering model plus a power-law, and we find a good fit in the “low” state with $\beta=0.49(1)$ , $kT=0.36$ (3), $\tau=1.2(1)$ and $\Gamma=1.0(1)$ (see Rea et al."
+ 2007b for details on the model)., 2007b for details on the model).
+" In the ""high"" state. & and AT donot vary. while we found ar =1.8(1) and Γ=0.8(I). which reflects the hardening of the spectrum."," In the “high” state, $\beta$ and $kT$ donot vary, while we found a $\tau=1.8(1)$ and $\Gamma=0.8(1)$, which reflects the hardening of the spectrum."
+" The ""high"" state best fit spectrum is shown in Fig. 2.."," The “high"" state best fit spectrum is shown in Fig. \ref{fig:spectrum}."
+ IBIS data alone do not allow a blind search for timing signatures., IBIS data alone do not allow a blind search for timing signatures.
+ In order to perform a timing analysis in the hard X-ray range. we first derived a phase coherent timing solution using publicly available RYTE//PCA data covering the same period of the IBIS data.," In order to perform a timing analysis in the hard X-ray range, we first derived a phase coherent timing solution using publicly available /PCA data covering the same period of the IBIS data."
+ We discovered two new glitches. close to the epoch of our IBIS observations 3 and 4. which are described in detail in ?..," We discovered two new glitches, close to the epoch of our IBIS observations 3 and 4, which are described in detail in \citet{israel07}. ."
+ We used the temporal solutions derived there (see Table 3)). for the time periods covering the ddata before and after the first new glitch (and prior the second one).," We used the temporal solutions derived there (see Table \ref{tab:time}) ), for the time periods covering the data before and after the first new glitch (and prior the second one)."
+ We then extracted the IBIS events in the 20-200 keV band from the pixels that were illuminated by the source for at least the of their surface., We then extracted the IBIS events in the 20–200 keV band from the pixels that were illuminated by the source for at least the of their surface.
+ To optimize the signal to noise ratio. we restricted our search to the high state periods. namely fall 2004 and spring 2005 (obs.33 and 4).," To optimize the signal to noise ratio, we restricted our search to the high state periods, namely fall 2004 and spring 2005 3 and 4)."
+ For 33 (obs.44 results are not reported. but are consistent) we folded the data and. by using the Z test we find that the signal is pulsed at à 5.7 level (chance probability of 1.6x 107%).," For 3 4 results are not reported, but are consistent) we folded the data and, by using the $Z^{2}$ test we find that the signal is pulsed at a $\sim$ $\sigma$ level (chance probability of $1.6\times10^{-8}$ )."
+ If we divide the data in two energy bands. 20-60 andkeV.. the significance lowers to 3.24 and c. respectively.," If we divide the data in two energy bands, 20–60 and, the significance lowers to 3.2 and $\sigma$ , respectively."
+ The pulse profile of the source. as a function of energy. is shown in reffig:corr..," The pulse profile of the source, as a function of energy, is shown in \\ref{fig:corr}."
+ As it can be seen. there is a clear energy dependence of the pulse shape morphology: the peak which is predominant in the soft band disappears and a secondary peak grows with energy becoming the most prominent one above ~8 keV. The pulsed fraction cannot be easily derived from the PCA (a no-imaging instrument) data due to the uncertainties in the background estimation.," As it can be seen, there is a clear energy dependence of the pulse shape morphology: the peak which is predominant in the soft band disappears and a secondary peak grows with energy becoming the most prominent one above $\sim$ 8 keV. The pulsed fraction cannot be easily derived from the PCA (a non-imaging instrument) data due to the uncertainties in the background estimation."
+ Besides. the IBIS data are not easy to handle. because of large background variations due the to fact that the source is at different off-axis angles during the differet pointings.," Besides, the IBIS data are not easy to handle, because of large background variations due the to fact that the source is at different off-axis angles during the different pointings."
+ In. addition. due to the large IBIS field of view. the events can be polluted by nearby variable sources.," In addition, due to the large IBIS field of view, the events can be polluted by nearby variable sources."
+ Nevertheless. by averaging the background over the entire IBIS observations. we could measure a pulsed fraction by dividing the difference between the maximum and the minimum of the folded light curves by the count rates derived from imaging.," Nevertheless, by averaging the background over the entire IBIS observations, we could measure a pulsed fraction by dividing the difference between the maximum and the minimum of the folded light curves by the count rates derived from imaging."
+ The resulting values are: ~25% (in the 20-60 keV band). ~60% keV)). and ~40% We did not find significant differences in the IBIS data às far às it concerns the pulsed profiles or the pulsed fraction between observations 34 and 4. i.e. before and after the third glitch reported in reffig:corr We also reanalyzed the archival PDS data taken with dduring the pointings listed in reftab:xray..," The resulting values are: $\sim$ (in the 20–60 keV band), $\sim$ ), and $\sim$ We did not find significant differences in the IBIS data as far as it concerns the pulsed profiles or the pulsed fraction between observations 3 and 4, i.e. before and after the third glitch reported in \\ref{fig:corr} We also reanalyzed the archival PDS data taken with during the pointings listed in \\ref{tab:xray}."
+ Using the timing solutionsderived. from. the simultaneous MECS data (22) we find. and report for the first time. pulsations in the PDS data. in the 20-200 keV energy band.," Using the timing solutionsderived from the simultaneous MECS data \citep{rea03,israel01}
+ we find, and report for the first time, pulsations in the PDS data, in the 20–200 keV energy band."
+ Using Z statistics we did not detect the pulsations in the first dataset. the shorter one. with high confidence c). while in the secod dataset they were detected at c.," Using $Z^{2}$ statistics we did not detect the pulsations in the first dataset, the shorter one, with high confidence $\sigma$ ), while in the second dataset they were detected at $\sigma$."
+ By computing the counts in excess of the phase minimum. we found a 20-200 keV pulsed flux of 0.02520.010 (for the second observation: the results for the first one are compatible but less significant).," By computing the counts in excess of the phase minimum, we found a 20–200 keV pulsed flux of $0.025
+\pm 0.010$ (for the second observation; the results for the first one are compatible but less significant)."
+ By assuming a power-law model with photon index D=| forthe pulsed spectrum (?).. this gives a flux of ~ 07x107! erg em s7!. slightly lower but still compatible with that derived from IBIS data by using thesame approach. namely ~107! erg em7 s7!.," By assuming a power-law model with photon index $\Gamma=1$ forthe pulsed spectrum \citep{kuiper06}, , this gives a flux of $\sim$ $0.7\times 10^{-11}$ erg $^{-2}$ $^{-1}$ , slightly lower but still compatible with that derived from IBIS data by using thesame approach, namely $\sim$$10^{-11}$ erg $^{-2}$ $^{-1}$ ."
+The dual ποιοι represeutatiou is casily obtained by a discrete Fourier transform of the former: letting P.=Ah. this is Observe finally that the dimeusiou of the Hilbert space is directly proportional o the mass AY of the particle.,"The dual momentum representation is easily obtained by a discrete Fourier transform of the former: letting $P_k = k h$, this is Observe finally that the dimension of the Hilbert space is directly proportional to the mass $M$ of the particle."
+ This is in line with our approach of performing he classical lit iu its more transparent plivsical form. by. keeping  to its real physical value aud by cousideriug a particle of larger aud larger mass A.," This is in line with our approach of performing the classical limit in its more transparent physical form, by keeping $\hbar$ to its real physical value and by considering a particle of larger and larger mass $M$."
+ By standard quantization procedures one therefore compute the matrix representations of the evolution operator., By standard quantization procedures one therefore compute the matrix representations of the evolution operator.
+ This has been effected in |5].., This has been effected in \cite{physd}.
+ We reproduce here he formulae holding in the position represeutation. where the state vector has components ο -- NV.," We reproduce here the formulae holding in the position representation, where the state vector has components $\psi(Q_j)$, $j=0,\ldots,N$ ."
+ The action of the free evolution (79° induced w the free rotation has matrix clemeuts The quantum Arnold cat evolution operator is the product Ut=ντ where (FP has been defined just above aud AX is the operator with matrix elements where 6.7 is the Kronecker delta. that corresponds to the impulsive part of the Hamiltonian.," The action of the free evolution $U^{free}$ induced by the free rotation $\frac{P^2}{2M}$ has matrix elements The quantum Arnol'd cat evolution operator is the product $U^{cat} = K U^{free}$, where $U^{free}$ has been defined just above and $K$ is the operator with matrix elements where $\delta_{k,l}$ is the Kronecker delta, that corresponds to the impulsive part of the Hamiltonian."
+ We now consider a 1iore complex svstem consisting of a singlelarge particle of mass AY aud of a umber £ of smaller particles of mass ii. that are also bound to move on the same vine.," We now consider a more complex system consisting of a singlelarge particle of mass $M$ and of a number $I$ of smaller particles of mass $m$, that are also bound to move on the same ring."
+ Let q; aud p;. /—l.....I the coordinates aud momenta of these particles. respectively.," Let $q_i$ and $p_i$, $i=1,\ldots,I$ the coordinates and momenta of these particles, respectively."
+ We require that also the pliase space of the simall particles is a two dimensional torus. of £=1 periodicity in the variables 4; aud of periodicity at in the momenta p;.," We require that also the phase space of the small particles is a two dimensional torus, of $L=1$ periodicity in the variables $q_i$ and of periodicity $\frac{mL}{T}$ in the momenta $p_i$."
+ We must therefore have where is an integer., We must therefore have where $n$ is an integer.
+ The wavefuuctious of the small particles also take on the forma of eqs. (8..9)).," The wave–functions of the small particles also take on the form of eqs. \ref{eq-ham5b}, \ref{eq-ham5c})),"
+ with» in place of jV., with $n$ in place of $N$ .
+ The manyparticles wavefunctious can be written ou the basis of ποιο eigen.fuuctious as, The many–particles wave–functions can be written on the basis of momentum eigen–functions as
+this object.,this object.
+ For instance. Stvohmaver(2001b). reported a 40 Iz QPO seen sometimes simultaneously with the more stable 67 Iz QPO (Morganοἱal.1997).," For instance, \citet{Strohmayer2001b} reported a 40 Hz QPO seen sometimes simultaneously with the more stable 67 Hz QPO \citep{Morgan1997}."
+. QPOs have been observed with (he following increasing frequencies 27-40-56-67 Iz (Bellonivanderlklis 2004).," QPOs have been observed with the following increasing frequencies 27-40-56-67 Hz \citep{Belloni2001,
+ vanderKlis2004}."
+. Some LF-QPOs varving between 1 and 15 Iz have been investigated intensivelv by Markwardtetal.(1999)., Some LF-QPOs varying between 1 and 15 Hz have been investigated intensively by \cite{Markwardt1999}.
+. Their properties seem to be different [rom the IIF-QPOs. probably implying a dillerent physical mechanism at work.," Their properties seem to be different from the HF-QPOs, probably implying a different physical mechanism at work."
+ We emphasize that ib is not the scope of this paper to explain all the observed QPOs but just to explain or predict the HF-QPOs., We emphasize that it is not the scope of this paper to explain all the observed QPOs but just to explain or predict the HF-QPOs.
+ LF-QPOs should be related to Lense-Thirring precession or to some other (magneto-)hyvdrodyvnamieal modes in the accretion disk., LF-QPOs should be related to Lense-Thirring precession or to some other (magneto-)hydrodynamical modes in the accretion disk.
+" Nevertheless. [rom the frequencies of the twin HF-QPOs. the constant angular pattern speed of the density wave is derived Irom our model by 74,=0/2214/215/3756 Iz."," Nevertheless, from the frequencies of the twin HF-QPOs, the constant angular pattern speed of the density wave is derived from our model by $\nu_{\rm w} = \Omega_{\rm w}/2\pi = \nu_1/2 = \nu_2/3 \approx 56$ Hz."
+ Pulling this value of the gravitational Ποια pattern speed into Table 1 we get the results shown in Table 4.., Putting this value of the gravitational field pattern speed into Table \ref{tab:Table} we get the results shown in Table \ref{tab:TableRes}.
+ There is a problem on selecting the relevant and meaningful frequencies in (his table., There is a problem on selecting the relevant and meaningful frequencies in this table.
+ Indeed. (he spiral density. wave travels outwards (o several tens of Schwarzschild radii.," Indeed, the spiral density wave travels outwards to several tens of Schwarzschild radii."
+ Áecording to Maoetal...(2003).. the precise extension of the propagation of the wave with constant Irequency ancl without attenuation depends on the accretion rate ancl on the viscosity. (at least for their axisvnunetric hvdrodynamical modes).," According to \citet{Mao2008}, the precise extension of the propagation of the wave with constant frequency and without attenuation depends on the accretion rate and on the viscosity (at least for their axisymmetric hydrodynamical modes)."
+ Assuming that the wave can travel only up to a radius roy. for a black hole of mass {μι and angular momentum “py the lowest excited Irequency is To give some orders of magnitude. let's use Vp=10M.capi0.raa30£4. we find My=19.7 Uz.," Assuming that the wave can travel only up to a radius $r_{\rm out}$ , for a black hole of mass $M_{\rm
+ BH}$ and angular momentum $a_{\rm BH}$, the lowest excited frequency is To give some orders of magnitude, let's use $M_{\rm BH} = 10 \,
+M_\odot, a_{\rm BH} = 0, r_{\rm out} = 30 \, R_{\rm g}$, we find $\nu_{\rm low} = 19.7$ Hz."
+ Which frequencies to select depends strongly on the spiral wave properties. how lar they can propagate. which azimuthal numbers 7? are excited. al which amplitude and so on.," Which frequencies to select depends strongly on the spiral wave properties, how far they can propagate, which azimuthal numbers $m$ are excited, at which amplitude and so on."
+ The detailed invesGeation is left for future work., The detailed investigation is left for future work.
+ We can oulv Claim that the observedQPOs are retrieved by our model., We can only claim that the observedQPOs are retrieved by our model.
+ How to select them remains, How to select them remains
+"More likely these differences are caused by crowding effects and undetected systematic errors of the different telescopes used in this investigation,",More likely these differences are caused by crowding effects and undetected systematic errors of the different telescopes used in this investigation.
+ MIO is a globular cluster with three slow variable stars (Clement et 1985) and several UV bright stars (Zinn et 1972: Harris et 1983)., M10 is a globular cluster with three slow variable stars (Clement et 1985) and several UV bright stars (Zinn et 1972; Harris et 1983).
+ Our sample contains only the variable stars V2 and V3 (see references in Clement et 1985)., Our sample contains only the variable stars V2 and V3 (see references in Clement et 1985).
+ The other special stars are located in crowded regions. below the limiting magnitude. or exhibit too large proper motion errors.," The other special stars are located in crowded regions, below the limiting magnitude, or exhibit too large proper motion errors."
+ Both variable stars V2. V3 have according to our investigation a membership probability of97%.," Both variable stars V2, V3 have according to our investigation a membership probability of."
+. This confirms the earlier suggestion (Clement et 1985) that these stars are of W-Virginis type. although V3 has a very unusual period of 7.8 days.," This confirms the earlier suggestion (Clement et 1985) that these stars are of W-Virginis type, although V3 has a very unusual period of 7.8 days."
+ The first absolute proper motion of MIO based on the Hipparcos system was given in Geffert et (1997). which was later used in Odenkirchen et (1997) and Dinescu et (1999).," The first absolute proper motion of M10 based on the Hipparcos system was given in Geffert et (1997), which was later used in Odenkirchen et (1997) and Dinescu et (1999)."
+ However. as mentioned in Odenkirchen et (1997). the result of M10 was only preliminary due to the small number of Hipparcos stars used for the reduction of the plates.," However, as mentioned in Odenkirchen et (1997), the result of M10 was only preliminary due to the small number of Hipparcos stars used for the reduction of the plates."
+ The use of additional plates from Shanghai has led to a significant lower proper motion in declination., The use of additional plates from Shanghai has led to a significant lower proper motion in declination.
+ However. even for this sample of plates the solution is only moderately stable: Omitting one reference star leads to an absolute proper motion of M10. which differs from the original solution by about more than the mean uncertainties of the absolute proper motions.," However, even for this sample of plates the solution is only moderately stable: Omitting one reference star leads to an absolute proper motion of M10, which differs from the original solution by about more than the mean uncertainties of the absolute proper motions."
+ Fortunately. this is not the case for the denser ACT catalogue.," Fortunately, this is not the case for the denser ACT catalogue."
+ In summary. the use of the ACT catalogue resulted in à more stable solution. while the use of Hipparcos stars provided the more direct link to an extragalactic reference frame.," In summary, the use of the ACT catalogue resulted in a more stable solution, while the use of Hipparcos stars provided the more direct link to an extragalactic reference frame."
+ The catalogues based on the ACT catalogue (Urban et 1998) and on the Hippareos catalogue (ESA 1997) contributed therefore with equal weight to the final absolute proper motion of the globular cluster M10., The catalogues based on the ACT catalogue (Urban et 1998) and on the Hipparcos catalogue (ESA 1997) contributed therefore with equal weight to the final absolute proper motion of the globular cluster M10.
+ We obtained a mean absolute proper motion of dcosó=5.51.2 mas/yr and ps=—6.241.2 mas/vr., We obtained a mean absolute proper motion of $\mua = -5.5 \plm 1.2\;$ mas/yr and $\mud = -6.2 \plm 1.2\;$ mas/yr.
+ Together with the distance from the Sun of 4.3 kpe and a radial velocity of |76 (Harris 1996) we have calculated the velocity components in a system of galactic standard at rest (L.V.W ). peri- and apogalactic distances Ry. Ry. eccentricity e and the s-component of the angular momentum of the orbit using a simple logarithmic galactic mass model (see e.g. Dauphole et 1996).," Together with the distance from the Sun of 4.3 kpc and a radial velocity of $+76\;$ (Harris 1996) we have calculated the velocity components in a system of galactic standard at rest $U$ $V$ $W$ ), peri– and apogalactic distances $R_{\rm p}$, $R_{\rm a}$, eccentricity $e$ and the $z$ -component of the angular momentum of the orbit using a simple logarithmic galactic mass model (see e.g. Dauphole et 1996)."
+ These data are shown in Table 7.., These data are shown in Table \ref{t_m10dat}.
+ Figure 2 gives the orbit of MIO integrated over 10 Gyr backwards according to the method of Allen Santillánn (1991) using the programme of Odenkirchen et (1997)., Figure \ref{f_m10orb} gives the orbit of M10 integrated over 10 Gyr backwards according to the method of Allen Santillánn (1991) using the programme of Odenkirchen et (1997).
+ Our results demonstrate the importance of the new astrometric catalogues (Hipparcos. ACT) for the determination of absolute proper motions of objects located in the galactic plane.," Our results demonstrate the importance of the new astrometric catalogues (Hipparcos, ACT) for the determination of absolute proper motions of objects located in the galactic plane."
+ For these objects no extragalactic link of the proper motion system is available due to the absense of suitable objects., For these objects no extragalactic link of the proper motion system is available due to the absense of suitable objects.
+ We show that the differences between the solutions based on the ACT catalogue and the corresponding one based on the Hipparcos catalogue are small., We show that the differences between the solutions based on the ACT catalogue and the corresponding one based on the Hipparcos catalogue are small.
+ Although this result has to be checked using data of other fields. it would offer a range of applications for the ACT catalogue because of its larger number of stars.," Although this result has to be checked using data of other fields, it would offer a range of applications for the ACT catalogue because of its larger number of stars."
+ Moreover. even reaching only moderate accuraey on the wide field CCD frames with respect to small field CCD observations. our results indicate that these observations may supersede the photographie plates in near future.," Moreover, even reaching only moderate accuracy on the wide field CCD frames with respect to small field CCD observations, our results indicate that these observations may supersede the photographic plates in near future."
+ The new data of MIO show an eccentricity of the orbit of 0.41. which would place the cluster now in the range of eccentricities from 0.4 to 0.8 found for the majority of the globular clusters (Odenkirchen et 1997).," The new data of M10 show an eccentricity of the orbit of 0.41, which would place the cluster now in the range of eccentricities from 0.4 to 0.8 found for the majority of the globular clusters (Odenkirchen et 1997)."
+ The rotational velocity changed to 9=|111417 wwith respect to the |144 ffound in Odenkirchen et al. (, The rotational velocity changed to $\Theta = +111 \plm 17\;$ with respect to the $+144\;$ found in Odenkirchen et al. (
+1997).,1997).
+ According to its moderate eccentricity and rotational velocity. MIO belongs rather to the halo class of objects.," According to its moderate eccentricity and rotational velocity, M10 belongs rather to the halo class of objects."
+ This ts in line with its metallicity., This is in line with its metallicity.
+ However. the z-distance from the galactic plane does not exceed 3 kpe. which would be more characteristic for a thick disk object.," However, the $z$ -distance from the galactic plane does not exceed 3 kpc, which would be more characteristic for a thick disk object."
+ou the order of 28 pc (Dahlia2008).,on the order of 28 pc \citep{Dahm}.
+. Tf we assume that the known members of NGC 2261 are biased. towards low extinctions. this could cause the distance we derive to the cluster to be slelth smaller than the true ecolctric center of the cluster.," If we assume that the known members of NGC 2264 are biased towards low extinctions, this could cause the distance we derive to the cluster to be slightly smaller than the true geometric center of the cluster."
+ The distance we derive here. however. would properly describe the distance of the population of currently known members: we do no consider this to be a systematic error ia our ajalvsis. bu rather a nuance that should inform the iuterpretation of our results.," The distance we derive here, however, would properly describe the distance of the population of currently known members; we do not consider this to be a systematic error in our analysis, but rather a nuance that should inform the interpretation of our results."
+ Finally. the fundamental assuuption of raidoni axia orientations is itself somewhat questionable since stellar clusters like NCC 2261 are predicted to have collapscc from siugle cloud complexes.," Finally, the fundamental assumption of random axial orientations is itself somewhat questionable since stellar clusters like NGC 2264 are predicted to have collapsed from single cloud complexes."
+ We might exect. then. for there to be a preferred oricutation of sars in the cluster resulting from the conservation of he cloud’s initial augular momenta. or possibly from t16 presence of large scale magnetic fields.," We might expect, then, for there to be a preferred orientation of stars in the cluster resulting from the conservation of the cloud's initial angular momentum, or possibly from the presence of large scale magnetic fields."
+ Previous appications of the siui distriution techuique have. however. produced results that are in agereenmient with precise parallax πιοολο (Preibisch&Smith1997:Jeffries2007).," Previous applications of the $\sin i$ distribution technique have, however, produced results that are in agreement with precise parallax measurements \citep{Preibisch, Jeffries}."
+. This aereenmeΕςi provides evidence iu support of the random axial ¢xieutaion assuniption. but it has not been clirectly coufirned by observations.," This agreement provides evidence in support of the random axial orientation assumption, but it has not been directly confirmed by observations."
+ The ISCUSSIOLL ayove has revealed a nuuber of poteutia SVSematic effects in our analysis., The discussion above has revealed a number of potential systematic effects in our analysis.
+ The five potential systenatics inost amenable το direct Investigation (the inclusion of CTTSs in our primary sample: the acopted iTuninositv. temperature. aud cesiu error distiibut1015: aid the asstumed binary fraction) all influence the «erived best fit distance at the ~ level.," The five potential systematics most amenable to direct investigation (the inclusion of CTTSs in our primary sample; the adopted luminosity, temperature, and $v \sin i$ error distributions; and the assumed binary fraction) all influence the derived best fit distance at the $\sim$ level."
+ Wetrervefore combine these individual uucertanties in (πιάταre to characterize the potential error in our best fit τιdistance due to systematic effects; procποιο a estimate of our total systematic uncertaiuties of x 12 or 110 pe.," We therefore combine these individual uncertanties in quadrature to characterize the potential error in our best fit distance due to systematic effects, producing a estimate of our total systematic uncertainties of $\pm$ 12, or 110 pc."
+ We have caleulated distance-dependent sin; values for a supe of OF prenuain sequence stars in the open cluster NGC 2261., We have calculated distance-dependent $\sin i$ values for a sample of 97 pre-main sequence stars in the open cluster NGC 2264.
+ By comparie the observed sins cistribuion to a modeled distribution assuming raudoni axial orientations. we derive a distiuice of 913 pe to NGC 2261: quantitative tests of our analvsis reveal sampling and svstemiatie errors of LO and 110 pe. respectivevy.," By comparing the observed $\sin i$ distribution to a modeled distribution assuming random axial orientations, we derive a distance of 913 pc to NGC 2264; quantitative tests of our analysis reveal sampling and systematic errors of 40 and 110 pc, respectively."
+ Our distance estimate does not rely on evolutionary noclels to any significant Our cistance estimate is sienificautlv higher han a munber of previously determined distances. particlarly the widely cited value of τοῦ pc found by Sunectal.(1997).," Our distance estimate does not rely on evolutionary models to any significant Our distance estimate is significantly higher than a number of previously determined distances, particularly the widely cited value of 760 pc found by \citet{Sung}."
+. Iu general. though. our estimate falls withiithe typical range of calculated distauc08 (730 pc to 950 pe) for NGC 2261.," In general, though, our estimate falls within the typical range of calculated distances (730 pc to 950 pc) for NGC 2264."
+" Table 2 provides a comparison o: OUI result to previous distance estiiates,", Table \ref{tab:prevdist} provides a comparison of our result to previous distance estimates.
+ A distance to NGC 2261 of 913 pc represents ani dcrcase of approximately compared to the widely execίσος value of τοῦ pe. though the two results are formally consistent withiithe stun of the statistical aud systemaic error bars.," A distance to NGC 2264 of 913 pc represents an increase of approximately compared to the widely excepted value of 760 pc, though the two results are formally consistent within the sum of the statistical and systematic error bars."
+" The mean age of NGC 2261 is commonly cited as ~ 3 Απο though there is evidence or a considerable age spread witlin the cluster (Dali2¢ΙΟ),"," The mean age of NGC 2264 is commonly cited as $\sim$ 3 Myrs, though there is evidence for a considerable age spread within the cluster \citep{Dahm}."
+" The huninosities of preanain sequence stars are often compared with oxedietious. of theoretical pre-miai) sequence models to infer the age of their parent «‘laster: as Dhuuudnostv declines through the preauain seqence phase. the larger ""uumositfies produced bv assuuii8 a ereater distance o the cluster will produce a votuser inferred age for he cluster."," The luminosities of pre-main sequence stars are often compared with predictions of theoretical pre-main sequence models to infer the age of their parent cluster; as luminosity declines through the pre-main sequence phase, the larger luminosities produced by assuming a greater distance to the cluster will produce a younger inferred age for the cluster."
+ Increasing the assunred distance to NGC 2261 from τοῦ to 910 pe changes the distance modulus w OL anas., Increasing the assumed distance to NGC 2264 from 760 to 910 pc changes the distance modulus by 0.4 mag.
+ The correspouding 0.lamnag brighteuiug of he stars produces a shift iu the :woe of the cluster., The corresponding 0.4 mag brightening of the stars produces a shift in the age of the cluster.
+" We wave produced a crude esuate of the revised age of the cluster bv determuning the age at which a 1 M. star's II baud magnitude is 0. mae brighter than at 3 \Vrs according to the preauai rsequence models calculated by Daraffeetal.(1998)... the distance we derive here uplies a revised age for NGC 2261 of ~1 5 Ας,"," We have produced a crude estimate of the revised age of the cluster by determining the age at which a 1 $_{\odot}$ star's H band magnitude is 0.4 mag brighter than at 3 Myrs: according to the pre-main sequence models calculated by \citet{Baraffe}, the distance we derive here implies a revised age for NGC 2264 of $\sim$ 1.5 Myrs."
+ The detailed analvsis of cluster members requireG0 for robust estimates of the age aud properties of NCC' 2261 πι light of our new derived distance. however. is bevond the scope of this work.," The detailed analysis of cluster members required for robust estimates of the age and properties of NGC 2264 in light of our new derived distance, however, is beyond the scope of this work."
+ We deteriuned the «istance to the open cluster NGC 2261 using a statistical analysis of cluster member iuchnations., We determined the distance to the open cluster NGC 2264 using a statistical analysis of cluster member inclinations.
+" We derived distance-depeudeut values of siu (where ¢ is the inclination :mele) for 97 stars in NGC 2261 from measired rotation periods. Iuninosities. effective tempcratures. and projected equatorial rotation velocities. esmi. of fjose stars,"," We derived distance-dependent values of $\sin i$ (where $i$ is the inclination angle) for 97 stars in NGC 2264 from measured rotation periods, luminosities, effective temperatures, and projected equatorial rotation velocities, $v \sin i$, of these stars."
+ We measured. 96 of the cesin/ values ini our sample by analyzing ligh-resolution spectra witQ1 a cross-correlation technique., We measured 96 of the $v \sin i$ values in our sample by analyzing high-resolution spectra with a cross-correlation technique.
+ We modeled the observed. distriπιjon of sin’ for the cluster by ΠΕ that meniy stars have random axial oricutatious aud by adoptinge prescriptions for the lüeasnrenient errors da our saldle., We modeled the observed distribution of $\sin i$ for the cluster by assuming that member stars have random axial orientations and by adopting prescriptions for the measurement errors in our sample.
+ By adjusting the distance asstuned in the observed VAa14 distribution uutil, By adjusting the distance assumed in the observed $\sin i$ distribution until
+ stars. as they expand (assuming virialised conditions).," stars, as they expand (assuming virialised conditions)."
+ Thus. if the planet formation timescale is Z5510? vears. stars in a cluster of 3«107 stars with densities Z310 wwill typically have their clises stripped. through stellar interactions before they are able to form gas giant. planets.," Thus, if the planet formation timescale is $\simgreat 5
+\times 10^6$ years, stars in a cluster of $3\times 10^5$ stars with densities $\simgreat 3 \times 10^{5}$ will typically have their discs stripped through stellar interactions before they are able to form gas giant planets."
+ Vhe οσοι of the number of stars on the necessary cluster conditions can be seen in Figure 3. which shows the critical densities for encounters within 10 and 50 (of a 1.0 sstar) on timescales o£ 2 and LO million vears as a function of the number of stars in the cluster., The effect of the number of stars on the necessary cluster conditions can be seen in Figure \ref{destNvel} which shows the critical densities for encounters within 10 and 50 (of a 1.0 star) on timescales of 2 and 10 million years as a function of the number of stars in the cluster.
+ Phe decrease in critical densities in lareer-N svstems is due to the larger velocity dispersions in such systems., The decrease in critical densities in larger-N systems is due to the larger velocity dispersions in such systems.
+" From this. it can be seen that densities 2107 start to become. interesting for disrupting some outer planets or the outer dise whereas densities 10"" are required to allect. planets that form at distances of =5AL..."," From this, it can be seen that densities $\simgreat 10^4$ start to become interesting for disrupting some outer planets or the outer disc whereas densities $\simgreat 10^5$ are required to affect planets that form at distances of $\approx 5$."
+ Combining the expected evolution of clusters of different numbers of stars and the expected: encounter timescale. we expect that encounters within LO can occur within the planet forming timescale in the large-N elobular clusters but not in the smaller-N open (embedded) clusters.," Combining the expected evolution of clusters of different numbers of stars and the expected encounter timescale, we expect that encounters within 10 can occur within the planet forming timescale in the large-N globular clusters but not in the smaller-N open (embedded) clusters."
+ This dilference is mainly due to the expected densities in the earliest stages of the eluster evolution combined with the higher critical densities for encounters in the low-N systems., This difference is mainly due to the expected densities in the earliest stages of the cluster evolution combined with the higher critical densities for encounters in the low-N systems.
+ Planctary systems that do succeed. in. forming in a stellar cluster are then subject to disruptions from stellar encounters., Planetary systems that do succeed in forming in a stellar cluster are then subject to disruptions from stellar encounters.
+ Once a planet has formed. and possibly migrated to its final separation from the parent star. the probability ofa clisruptive encounter depends on this separation and on the cluster properties.," Once a planet has formed, and possibly migrated to its final separation from the parent star, the probability of a disruptive encounter depends on this separation and on the cluster properties."
+ In order to quantify this probability we performed. simulations of the evolution of a population of planets in various cluster conditions., In order to quantify this probability we performed simulations of the evolution of a population of planets in various cluster conditions.
+ The initial distribution of planetary orbits was taken to be Hat in log separation. and spannec a range of 0.01 to 100Au.," The initial distribution of planetary orbits was taken to be flat in log separation, and spanned a range of 0.01 to 100."
+. Two main cluster types were investigated. elobular clusters and open clusters.," Two main cluster types were investigated, globular clusters and open clusters."
+ For the elobular cluster case. we chose a velocity. dispersion.: of⋅ LO.," For the globular cluster case, we chose a velocity dispersion of 10."
+ For+ the open cluster case. the velocity dispersion. is 2," For the open cluster case, the velocity dispersion is 2."
+ We then used Equation((1)) to caleulate the probability of encounters within various time intervals and cluster densities (sce Figures 4 and 5))., We then used \ref{eqtcoll}) ) to calculate the probability of encounters within various time intervals and cluster densities (see Figures \ref{case11} and \ref{case12}) ).
+ In cach case. the total time interval was split into smaller time steps. so that multiple encounters were possible.," In each case, the total time interval was split into smaller time steps, so that multiple encounters were possible."
+ The οσο of an encounter on a planetary. svsten was determined. as follows., The effect of an encounter on a planetary system was determined as follows.
+" Where the kinetic energy. of the perturber was greater than the binding energy. of the planet. the planetary system was assumed to be detroved (or ""jonised)."," Where the kinetic energy of the perturber was greater than the binding energy of the planet, the planetary system was assumed to be detroyed (or `ionised')."
+" Where the kinetic energy of the perturber was less than the binding energy of the planetary svstem. the planet was assumed to lose energy and move closer to its parent star (the svstem is ""hardened)."," Where the kinetic energy of the perturber was less than the binding energy of the planetary system, the planet was assumed to lose energy and move closer to its parent star (the system is `hardened')."
+ This hardening was taken to be25%... based on average values obtained in binarv-single star scattering experiments (Sigurdsson Phinney 1993: Davies 1997).," This hardening was taken to be, based on average values obtained in binary-single star scattering experiments (Sigurdsson Phinney 1993; Davies 1997)."
+" In the former case. the system is termed ‘soft’. while in the latter. it is said to be ""hard."," In the former case, the system is termed `soft', while in the latter, it is said to be `hard'."
+ The hard/soft boundary is given by where Aly.Ms.AL; are. the masses of the primary. secondary (in this case planet) and of the perturber star.," The hard/soft boundary is given by where $M_1, M_2, M_3$ are the masses of the primary, secondary (in this case planet) and of the perturber star."
+ The encounter velocity. (ge is essentially the velocity dispersion. Casp OL the cluster.," The encounter velocity, $\venc$ is essentially the velocity dispersion, , of the cluster."
+ From Equation 2. we see that most. planetary systems. where Ado is small are ‘soft’ and will easily be ionised through encounters.," From Equation \ref{hardsoft} we see that most planetary systems, where $M_2$ is small, are `soft' and will easily be ionised through encounters."
+ The question of hardenning is therefore not crucial to our conclusions., The question of hardenning is therefore not crucial to our conclusions.
+ ln addition. encounters can drive eccentricity into the planetary. system. ancl thus potentially further instability (Davies Sigurdsson 2000).," In addition, encounters can drive eccentricity into the planetary system, and thus potentially further instability (Davies Sigurdsson 2000)."
+ The results from our simulations are cividec into two sections depending on the chosen velocity dispersion., The results from our simulations are divided into two sections depending on the chosen velocity dispersion.
+ Firstly we present the results appropriate for Globular clusters with usp=10Lo then those for open clusters with (i2kms," Firstly we present the results appropriate for Globular clusters with $\vdsp =10$, then those for open clusters with $\vdsp = 2$ ."
+ Figure 4. presents the results for the planetary systems in clusters where the velocity dispersion is z10|., Figure \ref{case11} presents the results for the planetary systems in clusters where the velocity dispersion is $\approx 10$.
+.. The figure is broken up into dillerent panels each appropriate for a specific stellar density and for a specific amount of time., The figure is broken up into different panels each appropriate for a specific stellar density and for a specific amount of time.
+ Globular' clusters have typical. core densities- of⋅ =I01. mean densities of zz107 aand lifetimes that extend to the age ofthe Galaxy (I0σης.104? vears).," Globular clusters have typical core densities of $\approx 10^4$, mean densities of $\approx 10^3$ and lifetimes that extend to the age of the Galaxy $10^9 \simless t \simless 10^{10}$ years)."
+ The cluster densities increase from left to right and the time increases from top to bottom., The cluster densities increase from left to right and the time increases from top to bottom.
+ The dotted lines eive the initial planetary. distribution while the heavy solid lines give the final planetary. distribution., The dotted lines give the initial planetary distribution while the heavy solid lines give the final planetary distribution.
+ The light. solid lines indicate the fraction of planetary svstems. ancl their initial separations. that have been ionisect.," The light solid lines indicate the fraction of planetary systems, and their initial separations, that have been ionised."
+ The overall result. is that the wider. (softer) systems are more casily disrupted. than are the tighter syvstenis., The overall result is that the wider (softer) systems are more easily disrupted than are the tighter systems.
+ Their larger cross section. for an encounter results. in lower critical densities and in shorter encounter timescales., Their larger cross section for an encounter results in lower critical densities and in shorter encounter timescales.
+ ‘Tighter systems require higher stellar densities in order to ensure a reasonable timescale for an encounter., Tighter systems require higher stellar densities in order to ensure a reasonable timescale for an encounter.
+ Furthermore. a cluster of given stellar density will disrupt. increasingly tighter systems with time until reaching the haresoft boundary.," Furthermore, a cluster of given stellar density will disrupt increasingly tighter systems with time until reaching the hard/soft boundary."
+ Thus. wide svstenis are disrupted in most clusters although they do require longer timescales in the least dense clusters.," Thus, wide systems are disrupted in most clusters although they do require longer timescales in the least dense clusters."
+ In contrast. the tighter systems are only disrupted in sullicientlv old. dense. clusters.," In contrast, the tighter systems are only disrupted in sufficiently old, dense, clusters."
+" It can be seen that in the case of eai,=LOla all svstems are soft."," It can be seen that in the case of $\vdsp = 10$, all systems are soft."
+ There is no hardening (moving planets to smaller separations) and any system. that is perturbed isdisrupted., There is no hardening (moving planets to smaller separations) and any system that is perturbed isdisrupted.
+ This occurs as the hard/soft boundary. is at Ryo80.02. or 4 (smaller than the separations, This occurs as the hard/soft boundary is at $\rhd \approx 0.02$ or 4 (smaller than the separations
+that form locally can then hierarchically merge to form large stellar clusters (??)..,"that form locally can then hierarchically merge to form large stellar clusters \citep{BonBatVin2003, Bate2009a}."
+ Residual gas in these bound regions then falls into the gravitational potential of the cluster to be competitively accreted by the growing massive stars located in the bottom of the potential well (??)..," Residual gas in these bound regions then falls into the gravitational potential of the cluster to be competitively accreted by the growing massive stars located in the bottom of the potential well \citep{BonVinBat2004,BonBat2006}."
+ The massive stars are thus located in the stellar clusters., The massive stars are thus located in the stellar clusters.
+ The majority of the stars and brown dwarfs formed are in high-density regions or have been ejected from stellar clusters through interactions (??)..," The majority of the stars and brown dwarfs formed are in high-density regions or have been ejected from stellar clusters through interactions \citep{BatBonBro2002a,BatBonBro2003}."
+" As noted in (?),, the brown dwarfs predominantly form in stellar clusters due to the compression of the gas to high local densities as it falls into the gravitational potential."," As noted in \citep{BonClaBat2008}, the brown dwarfs predominantly form in stellar clusters due to the compression of the gas to high local densities as it falls into the gravitational potential."
+ The high-mass stars are also predominantly formed in clusters (??)..," The high-mass stars are also predominantly formed in clusters \citep{BonVinBat2004, SmiLonBon2009}."
+ This leads to a potentially observable difference in the stellar IMFs of distributed and clustered star formation., This leads to a potentially observable difference in the stellar IMFs of distributed and clustered star formation.
+ Figure[] shows the final IMF for the overall population and also for distributed and clustered populations defined as those with a stellar density lower or higher than 100 stars pc?., Figure \ref{imfdenscut} shows the final IMF for the overall population and also for distributed and clustered populations defined as those with a stellar density lower or higher than 100 stars $^{-3}$.
+ The right-hand panel of figure] shows the corresponding populations where they are separated by their maximum stellar density during their evolution.," The right-hand panel of figure \ref{imfdenscut}
+ shows the corresponding populations where they are separated by their maximum stellar density during their evolution."
+ The cumulative distributions (figure [6)) show that the two distributions are statistically different and inconsistent (at the 1x10” level) with being drawn from the same population., The cumulative distributions (figure \ref{cimfdenscut}) ) show that the two distributions are statistically different and inconsistent (at the $1\times 10^{-13}$ level) with being drawn from the same population.
+ The distributed12 population has a significantly higher median stellar mass and a pronounced lack of low-mass objects., The distributed population has a significantly higher median stellar mass and a pronounced lack of low-mass objects.
+ This result helps explain the seemingly anomolous IMF in Taurus which appears to have a lack of brown dwarfs and high-mass stars in a distributed population (?).., This result helps explain the seemingly anomolous IMF in Taurus which appears to have a lack of brown dwarfs and high-mass stars in a distributed population \citep{Luhman2004b}.
+ One of the central questions we wish to address in this paper is the relationship between the nature and efficiency of the star formation process., One of the central questions we wish to address in this paper is the relationship between the nature and efficiency of the star formation process.
+ Previous studies (ie.??7?) showed that unbound clouds resulted in inefficient star formation.," Previous studies \citep[i.e.][]{ClaBon2004, Clarketal2005, ClaBonKle2008} showed that unbound clouds resulted in inefficient star formation."
+" Furthermore, the efficiency reduces dramatically the further the clouds are from being bound."," Furthermore, the efficiency reduces dramatically the further the clouds are from being bound."
+ Turbulent compression and shocks results in some star formation in these clouds but it is localised and much of the cloud escapes without entering the star formation process., Turbulent compression and shocks results in some star formation in these clouds but it is localised and much of the cloud escapes without entering the star formation process.
+" In the present study, we have one cloud that has regions which are bound and regions which are unbound with a spatially varying M/L from 0.6 to 1.4."," In the present study, we have one cloud that has regions which are bound and regions which are unbound with a spatially varying $M/L$ from 0.6 to 1.4."
+ This results in a range in local star formation efficiencies from 0.006 to 0.4., This results in a range in local star formation efficiencies from 0.006 to 0.4.
+ Figure[/] plots the local star formation efficicency as a function of the local binding of the cloud in terms of the critical M/L for the cloud to be globally bound., Figure \ref{sfevsmol} plots the local star formation efficicency as a function of the local binding of the cloud in terms of the critical $M/L$ for the cloud to be globally bound.
+ We see that after 1.02 free- times or ©6.6x10° years (and z3.9x10? years after the first stars formed) the local efficiency of star formation is strongly dependent on local binding of the cloud., We see that after $1.02$ free-fall times or $\approx 6.6 \times 10^5$ years (and $\approx 3.9\times 10^5$ years after the first stars formed) the local efficiency of star formation is strongly dependent on local binding of the cloud.
+ The bound, The bound
+The hydrogen in the Universe became predominantly neutral at the epoch of recombination. z;=1070.,"The hydrogen in the Universe became predominantly neutral at the epoch of recombination, $z=1070$."
+ The inter-galactic medium (IGM) out to z=5. however. 1s highly tonised.," The inter-galactic medium (IGM) out to $z=5$, however, is highly ionised."
+ Two observations made within the last decade have narrowed the window within which cosmic hydrogen retonisation took place., Two observations made within the last decade have narrowed the window within which cosmic hydrogen reionisation took place.
+ First. the electron scattering optical depth to the cosmic microwave background measured by the implies a redshift of reionisation of 114+1.4 (?)..," First, the electron scattering optical depth to the cosmic microwave background measured by the implies a redshift of reionisation of $11.4\pm 1.4$ \citep{Dunkley_etal:2009}."
+ Second. observations of the Lya forest in the spectra of z~6 quasars Indicate that ς=5.8 marks the tatl-end of the epoch of retonisation (?)..," Second, observations of the $\alpha$ forest in the spectra of $z\sim6$ quasars indicate that $z=5.8$ marks the tail-end of the epoch of reionisation \citep{Fan_etal:2006}."
+ These two results suggest that reionisation is an extended process., These two results suggest that reionisation is an extended process.
+ Therefore. in order to understand the chronology of reionisation in detail. there is considerable interest in detecting sources beyond z=6.4. the redshift of the most distant quasars so far discovered (??)..," Therefore, in order to understand the chronology of reionisation in detail, there is considerable interest in detecting sources beyond $z=6.4$, the redshift of the most distant quasars so far discovered \citep{Fan:2003, Willott:2007}."
+ Lye emitting galaxies (LAEs) have been detected out to <= 7(??)..," $\alpha$ emitting galaxies (LAEs) have been detected out to $z=7$ \citep{Kashikawa:2006, Ota_etal:2008}."
+ Since the strength of the Lyc emission line depends on the neutral fraction. vy). one approach advocated to chart the progress of retonisation (2). 1s to measure the evolution of the properties (e.g. abundance. clustering. ete) of LAEs.," Since the strength of the $\alpha$ emission line depends on the neutral fraction, $\xhi$, one approach advocated to chart the progress of reionisation \citep{McQuinn:2007} is to measure the evolution of the properties (e.g. abundance, clustering, etc) of LAEs."
+ However. the processes involved are difficult to model accurately (e.g. ??).. and there is little consensus vet on the interpretation of the results from this approach.," However, the processes involved are difficult to model accurately \citep[e.g.][]{Tasitsiomi:2006, Dijkstra:2007}, and there is little consensus yet on the interpretation of the results from this approach."
+ The confirmation of the cosmological nature of gamma ray bursts (GRBs) (?) revealed their potential as probes of the high-redshift Universe., The confirmation of the cosmological nature of gamma ray bursts (GRBs) \citep{Kulkarni_etal:1998} revealed their potential as probes of the high-redshift Universe.
+ ? emphasise two particular advantages of GRBs over quasars for studying reionisaton: that their afterglows will be visible to redshifts of z~10: and that the size of the ionised region around the host galaxy. which complicates the interpretation of the spectrum. will be small.," \cite{Barkana:2004} emphasise two particular advantages of GRBs over quasars for studying reionisaton: that their afterglows will be visible to redshifts of $z\sim10$; and that the size of the ionised region around the host galaxy, which complicates the interpretation of the spectrum, will be small."
+ Recent numerical simulations by ? and ?.. however. suggest that GRB host galaxies may be located in dense environments. and therefore the IGM immediately surrounding the host galaxy may be tonised by nearby quasars or local massive star forming galaxies. and so the size of the ionised region needs to be included in the modelling.," Recent numerical simulations by \cite{McQuinn_etal:2008} and \cite{Mesinger_Furlanetto:2008}, however, suggest that GRB host galaxies may be located in dense environments, and therefore the IGM immediately surrounding the host galaxy may be ionised by nearby quasars or local massive star forming galaxies, and so the size of the ionised region needs to be included in the modelling."
+ In any case a high signal-to-noise ratio (S/N) spectrum is needed to distinguish between the different signatures of a high-column density of neutral hydrogen in the GRB host galaxy (or nearby) (?)..D and of distributed neutral hydrogen in the IGM.," In any case a high signal-to-noise ratio (S/N) spectrum is needed to distinguish between the different signatures of a high-column density of neutral hydrogen in the GRB host galaxy (or nearby) \citep{RuizVelasco_etal:2007}, and of distributed neutral hydrogen in the IGM."
+ Until the launch of theSwift satellite (?).. the furthest known GRB was at z=4.5 (?).. compared to z=5.8 for the furthest known quasar at that time (2)..," Until the launch of the satellite \citep{Gehrels_etal:2004}, the furthest known GRB was at $z=4.5$ \citep{Andersen_etal:2000}, , compared to $z=5.8$ for the furthest known quasar at that time \citep{Fan_etal:2000}."
+ Since then. three GRBs at z>6 have been discovered usingSwifr: GRB 050904 at z=6.3 in 2005 (22); GRB 080913 at z26.7 in 2008 (?):: and. in 2009. the remarkable source GRB 090423 at z=82 (??).. the most distant object yet found.," Since then, three GRBs at $z>6$ have been discovered using: GRB 050904 at $z=6.3$ in 2005 \citep{Cusumano_etal:2006, Kawai_etal:2006}; GRB 080913 at $z=6.7$ in 2008 \citep{Greiner_etal:2009}; and, in 2009, the remarkable source GRB 090423 at $z=8.2$ \citep{Tanvir_etal:2009,Salvaterra_etal:2009}, the most distant object yet found."
+ ? present a detailed analysis of the optical afterglow spectrum. of GRB 050904. which was measured with the highest S/N of the three z>6 GRBs.," \cite{Totani_etal:2006} present a detailed analysis of the optical afterglow spectrum of GRB 050904, which was measured with the highest S/N of the three $z>6$ GRBs."
+ The spectrum displays ἃ red damping wing from Lye absorption that is produced by some combination of à high-column density absorber near the GRB (hereafter referred to as a DLA. for damped Lyw absorber). and a smoothly distributed component in the IGM.," The spectrum displays a red damping wing from $\alpha$ absorption that is produced by some combination of a high-column density absorber near the GRB (hereafter referred to as a DLA, for damped $\alpha$ absorber), and a smoothly distributed component in the IGM."
+ By fitting absorption models containing these two components. ? were able to place a limit on the neutral fraction of the IGM at z=6.3 of vy)«0.17(0.60) at 68%(95%) confidence.," By fitting absorption models containing these two components, \cite{Totani_etal:2006} were able to place a limit on the neutral fraction of the IGM at $z=6.3$ of $\xhi<0.17(0.60)$ at $68\%(95\%)$ confidence."
+ The constraints are relatively weak despite the reasonably high S/N of the spectrum., The constraints are relatively weak despite the reasonably high S/N of the spectrum.
+ The sources GRB 080913 and GRB 090423 are potentially more interesting. because of their higher redshifts. but the published spectra have insufficient. S/N to place any useful constraints on vy; when employing two-component (DLA-IGM) fits.," The sources GRB 080913 and GRB 090423 are potentially more interesting, because of their higher redshifts, but the published spectra have insufficient S/N to place any useful constraints on $\xhi$ when employing two-component (DLA+IGM) fits."
+ Here we deseribe and analyse an improved spectrum of GRB 080913. which includes unpublished spectroscopic data taken three nights after the published spectrum.," Here we describe and analyse an improved spectrum of GRB 080913, which includes unpublished spectroscopic data taken three nights after the published spectrum."
+ In Section 2 we describe the observations taken on each night. and the reduction techniques.," In Section 2 we describe the observations taken on each night, and the reduction techniques."
+ We analyse the combined spectrum in Section 3. first searching for absorption lines in the new spectrum. in order to measure the redshift of the source. and then fitting a two-component DLA+IGM model to the observed continuum break.," We analyse the combined spectrum in Section 3, first searching for absorption lines in the new spectrum, in order to measure the redshift of the source, and then fitting a two-component DLA+IGM model to the observed continuum break."
+Finally. the results are summarised in Section 4.,"Finally, the results are summarised in Section 4."
+ ? (hereafter Paper D) provide a summary of all the photometric and spectroscopic observations of GRB 080913., \cite{Greiner_etal:2009} (hereafter Paper I) provide a summary of all the photometric and spectroscopic observations of GRB 080913.
+ Here we recapthe details of the spectroscopic observations only., Here we recapthe details of the spectroscopic observations only.
+ The object, The object
+We present below results from previous studies of the hard ray spectrum of the sources in the sample.,We present below results from previous studies of the hard X--ray spectrum of the sources in the sample.
+ NGC 4151 is a nearby Seyfert 1.5 galaxy. and the brightest persistent AGN in the kkeV. energy band (after the radio-loud blazar Cen A).," NGC 4151 is a nearby Seyfert 1.5 galaxy, and the brightest persistent AGN in the keV energy band (after the radio-loud blazar Cen A)."
+ It hosts a BH with a mass of 4.678%10Μ... as estimated from reverberation mapping studies (Bentzetal..2006).. and radiates at a rate of of Leo/Lupp=0.014 (Crenshaw&Kraemer.2007).," It hosts a BH with a mass of $4.6^{+0.6}_{-0.5}{\times}10^{7}\,{\rm M}_{\odot}$, as estimated from reverberation mapping studies \citep{bentz06}, and radiates at a rate of of ${\rm L}_{\rm Bol}/{\rm L}_{\rm EDD}=0.014$ \citep{crenshaw07}."
+. Its spectrum has been extensively studied from. radio wavelengths to hard X-rays., Its spectrum has been extensively studied from radio wavelengths to hard X-rays.
+ Its energy spectrum above 20 keV is well fitted by a power-law model (PL. hereafter). with an exponential cut-off above energies 50-200 keV (see e.g. Piro et al..," Its energy spectrum above 20 keV is well fitted by a power-law model (PL, hereafter), with an exponential cut-off above energies 50–200 keV (see e.g. Piro et al.,"
+ 1999 and 2000: Beckmann et al..," 1999 and 2000; Beckmann et al.,"
+ 2005 and 2009; Lubinsski et al., 2005 and 2009; Lubińsski et al.
+.2010).,",2010)."
+ Lubinsski et al. (, Lubińsski et al. (
+"2010) in. particular. have suggested that the cut-off energy. E.. anti-correlates with the source flux. increasing from ~5—80 keV. at high fluxes. to ~100—200 keV. when the source tis at ""dim flux state.","2010) in particular, have suggested that the cut-off energy, $E_{\rm c}$, anti-correlates with the source flux, increasing from $\sim 5-80$ keV, at high fluxes, to $\sim 100-200$ keV, when the source is at “dim flux state""."
+ The spectral slope of the PL component. F. is rather hard (compared to other. X-ray bright. Seyfert galaxies). with best-fit values being ~1.4—1.6 (see e.g. Zdziarski et al..," The spectral slope of the PL component, $\Gamma$, is rather hard (compared to other, X–ray bright, Seyfert galaxies), with best-fit values being $\sim 1.4-1.6$ (see e.g. Zdziarski et al.,"
+ 1996: Petrueci et al..," 1996; Petrucci et al.,"
+ 2001: Winter et al..," 2001; Winter et al.,"
+ 2009)., 2009).
+ de Rosa et al. (, de Rosa et al. (
+2007) have suggested that there may exist flux related. intrinsic spectral slope variations in this source.,"2007) have suggested that there may exist flux related, intrinsic spectral slope variations in this source."
+" In addition to the PL component. observations first established the presence of a ""reflection hump” at energies above 10 keV (Zdziarski et al.."," In addition to the PL component, observations first established the presence of a “reflection hump"" at energies above 10 keV (Zdziarski et al.,"
+ 1996)., 1996).
+ This component has been detected by all major hard X-ray satellites since then. but with a variable amplitude. R. which ranges between 0.01 and almost | (re. the value expected in the case of à point source located on top of the aceretion dise. at relatively large height).," This component has been detected by all major hard X–ray satellites since then, but with a variable amplitude, $R$, which ranges between 0.01 and almost 1 (i.e. the value expected in the case of a point source located on top of the accretion disc at relatively large height)."
+ Finally. it is well established that the X-ray spectrum of NGC 4151 is also affected by the presence of neutral material. perhaps with a variable covering factor. and a column density. Ny. less than ~107 em7 (see e.g. Lubirisski et al..," Finally, it is well established that the X-ray spectrum of NGC 4151 is also affected by the presence of neutral material, perhaps with a variable covering factor, and a column density, $_{\rm H}$, less than $\sim 10^{23}$ $^{-2}$ (see e.g. Lubińsski et al.,"
+ 2010. and references therein).," 2010, and references therein)."
+ NGC 4945 is à nearby MMpe: Mauersberger et al.1996) Seyfert 2 galaxy. with a BH mass estimate of x135x109M. (Greenhilletal..1997).. which radiates at a rate of Lg/Lipp=0.10—0.60 (Madejskietal..2000).," NGC 4945 is a nearby Mpc; Mauersberger et al.1996) Seyfert 2 galaxy, with a BH mass estimate of ${\approx}1.4{\times}10^{6}\,{\rm M}_{\odot}$ \citep{greenhill97}, which radiates at a rate of ${\rm L}_{\rm Bol}/{\rm L}_{\rm EDD}=0.10-0.60$ \citep{madejski00}."
+. observations revealed the presence of a hard X-ray source in this galaxy. which is heavily absorbed at energies below 10 keV by neutral material with an Ny as high as ~1077 (Madejskietal..2000;," observations revealed the presence of a hard X–ray source in this galaxy, which is heavily absorbed at energies below 10 keV by neutral material with an $_{\rm H}$ as high as $\sim 10^{24}$ \citep {madejski00,guainazzi00,takeshi08}."
+An alternative method for removing location paraueter biases is fo use aji approach that is based on selective comparisons.,An alternative method for removing location parameter biases is to use an approach that is based on selective comparisons.
+ The SOFT method is founded on the asstnption that two objects with very similar physical pariuneters wwill necessarily have verv siwilar light curve shapes., The SOFT method is founded on the assumption that two objects with very similar physical parameters will necessarily have very similar light curve shapes.
+" The biases that we are trviug to remove arise primarily due to light curve shape nisumatclies between the template and the caididate SN,", The biases that we are trying to remove arise primarily due to light curve shape mismatches between the template and the candidate SN.
+ Thus. if we could ideutifv à eroup of SNe that have very simular vvalues. then we would expect a direct conparisou between eroup leubers to be largely free of bias.," Thus, if we could identify a group of SNe that have very similar values, then we would expect a direct comparison between group members to be largely free of bias."
+ Oue possibility would be to select candidate SNe that have vers* close iiatclies in the template library., One possibility would be to select candidate SNe that have very close matches in the template library.
+" For cxampο, we could pick out oulv those candidate objects that have a very strong match with a single template by elevatiug our template selection threshold to >90% or 95©."," For example, we could pick out only those candidate objects that have a very strong match with a single template by elevating our template selection threshold to $\gamma>90\%$ or $95\%$."
+ However. with oulv a ew dozen templates in the library most SN ciudidates will not lave a match with any template. so the available siuuple size would be severely. reduced.," However, with only a few dozen templates in the library most SN candidates will not have a match with any template, so the available sample size would be severely reduced."
+ A modification of this approach could overcome the hnütatious of a sparse template library bx leveraging the τον huge SN data sets to come from). future surveys., A modification of this approach could overcome the limitations of a sparse template library by leveraging the very large SN data sets to come from future surveys.
+ Suppose that a candidate SN X is observed at an unknown location 0x and we compare it against our N templates to eot a vector of integrated imenmiboership erades Ex=(xpovvNNS," Suppose that a candidate SN X is observed at an unknown location $_X$ and we compare it against our N templates to get a vector of integrated membership grades ${\bf
+ \Gamma_X}=(\gamma_{X1},\gamma_{X2},\ldots \gamma_{XN})$."
+ A second candidate SN Y exists at a differen location Q4- aud has a different set of membership grades Py-., A second candidate SN Y exists at a different location $_Y$ and has a different set of membership grades ${\bf \Gamma}_Y$.
+ If hese two objects have very. similar physical characteristics. Py=By. then we should expect that the same templates that are good natches for SN X. will also be similar to SN Y. so tha TyzTy.," If these two objects have very similar physical characteristics, $\bfPhi_X\approx\bfPhi_Y$, then we should expect that the same templates that are good matches for SN X will also be similar to SN Y, so that ${\bf \Gamma}_X \approx {\bf \Gamma}_Y$."
+" This provides a straightforward means for collecting a group of simular SNe. which we will call a “peer group."""," This provides a straightforward means for collecting a group of similar SNe, which we will call a “peer group.”"
+ We compare all the available candidates against the template models. and then eroup together candidates that have a very simular pattern of membership grades P.," We compare all the available candidates against the template models, and then group together candidates that have a very similar pattern of membership grades ${\bf \Gamma}$."
+ The members of this peer group generally will not have a perfect match in the template library. so no single template can provide an unbiased estimate for auv of the eroup members.," The members of this peer group generally will not have a perfect match in the template library, so no single template can provide an unbiased estimate for any of the group members."
+ However. we do expect that any meniber of this peer eroup could provide an unbiased estimate for the location parameters of any other member in the exoup.," However, we do expect that any member of this peer group could provide an unbiased estimate for the location parameters of any other member in the group."
+ All that is needed is to redefine the candidate SN members to take he place of the template library., All that is needed is to redefine the candidate SN members to take the place of the template library.
+ Recognizing that absolute distances are not recessary for cosmology. let us reduce the location vectors for our two peer SNe X and Y down ο a pair of threc-dinieusional sub-vectors ὃν=(τν.Atohaud My=(ry.yey.Faby).," Recognizing that absolute distances are not necessary for cosmology, let us reduce the location vectors for our two peer SNe X and Y down to a pair of three-dimensional sub-vectors $\bfvartheta_X=(z_X,A_{V,X},t_{pk,X})$ and $\bfvartheta_Y=(z_Y,A_{V,Y},t_{pk,Y})$."
+ The xocedure for executing a peer to peer Colparisol is as follows: At the end of this process we will lave a pair of rest-frame spectrophotometiic models fy(A.f) and fy(A.t) eiviug the flux as a function of wavelength aud time for both SN X and SN Y. at the assumed locations Jy and Oy.," The procedure for executing a peer to peer comparison is as follows: At the end of this process we will have a pair of rest-frame spectrophotometric models $f_X(\lambda,t)$ and $f_Y(\lambda,t)$ giving the flux as a function of wavelength and time for both SN X and SN Y, at the assumed locations $\bfvartheta_X$ and $\bfvartheta_Y$."
+ We now need to consider the relative Iuuinosity distances for these two objects., We now need to consider the relative luminosity distances for these two objects.
+ Assmuiug that membership in the same peer eroup indicates that these two objects are physically almost identical. then they should have nearly the same intrinsic hDnuuinosity.," Assuming that membership in the same peer group indicates that these two objects are physically almost identical, then they should have nearly the same intrinsic luminosity."
+ This mcaus that their relative Iuninosity distances Dyx and Dyy ave eucoded in the flux ratio:," This means that their relative luminosity distances $D_{L,X}$ and $D_{L,Y}$ are encoded in the flux ratio:"
+2=1.3.,$z=1.3$.
+ The clear break in the spectrum implies that a fairly accurate redshift is obtainable., The clear break in the spectrum implies that a fairly accurate redshift is obtainable.
+ In general. given the elliciency of STIS ancl the spectral shape of the afterglow. the S/N increases with higher redshift (see lig. 2]])," In general, given the efficiency of STIS and the spectral shape of the afterglow, the S/N increases with higher redshift (see fig. \ref{fig:int}] ])"
+ for the NUW ALAALA detector., for the NUV MAMA detector.
+ H£ the burst occurs at a redshift 21.5 then. with only a few orbits. one could determine the position of the Lyman break well enough to determine its redshift to an accuracy of better than 5 percent: it may even be possible to detect a a forest at high z with only a few orbits.," If the burst occurs at a redshift $z \age 1.5$ then, with only a few orbits, one could determine the position of the Lyman break well enough to determine its redshift to an accuracy of better than 5 percent; it may even be possible to detect a $\alpha$ forest at high $z$ with only a few orbits."
+ As it is easier to detect a redshift for bursts that originate from a higher z with the NUV. ALAALA detector. the absence of a Lyman limit in the spectrum would. in general. always provide an upper limit to the redshift.," As it is easier to detect a redshift for bursts that originate from a higher $z$ with the NUV MAMA detector, the absence of a Lyman limit in the spectrum would, in general, always provide an upper limit to the redshift."
+ From figure (2)). it is clear that use of the ALAALA detectors onboard the LIST. given. their very low internal countrates. are preferable to the CCD.," From figure \ref{fig:int}) ), it is clear that use of the MAMA detectors onboard the HST, given their very low internal countrates, are preferable to the CCD."
+ The disadvantage of using LST to infer a redshift is the αποεν of altering its scheduled: observations even on the timescale of weeks., The disadvantage of using HST to infer a redshift is the difficulty of altering its scheduled observations even on the timescale of weeks.
+ Certainly the advantage is that there are very. lew available grouncd-basecl telescopes that can resolve Ale LL and IV. absorption lines from faint objects: even then. such telescopes may not have the proper viewing conditions to the source and detected. lines may place dilferent limits in the redshift.," Certainly the advantage is that there are very few available ground-based telescopes that can resolve Mg II and C IV absorption lines from faint objects; even then, such telescopes may not have the proper viewing conditions to the source and detected lines may place different limits in the redshift."
+ Lo addition. there may be important emission (e.g. lines) in the UV. spectrum of GRBs that would not be observable from the ground.," In addition, there may be important emission (e.g. lines) in the UV spectrum of GRBs that would not be observable from the ground."
+ As 1e ensemble of optical counterparts begins to grow in size. it will be. possibl| to infer both their. recshilt distribution and the distribution of GRBs with respect to observable galaxies by noting the frequeney of absorption lines ancl detected Lyman limit (see fie. 1)).," As the ensemble of optical counterparts begins to grow in size, it will be possible to infer both their redshift distribution and the distribution of GRBs with respect to observable galaxies by noting the frequency of absorption lines and detected Lyman limit (see fig. \ref{fig:lls}) )."
+ ΙΓ no Lyman limit is detected in the spectra of the optical transients (especially at low-redshifts). the conclusion would be that either GRBs do not originate near galaxies (< 300) kpe or they are Galactic in origin.," If no Lyman limit is detected in the spectra of the optical transients (especially at low-redshifts), the conclusion would be that either GRBs do not originate near galaxies $\ale 300$ ) kpc or they are Galactic in origin."
+ A precise redshift’ of a gamma-ray burst would ereathy further the field. by providing an accurate understanding of the energies involved. in the burst., A precise redshift of a gamma-ray burst would greatly further the field by providing an accurate understanding of the energies involved in the burst.
+ Assuming the extrapolation of the spectral index (0) from the optical to UV. passhancl is correct. we find that a ALAALA detector observation of the bright afterglow of a GRB by STIS over 1 orbit could reveal a Lyman limit in the spectrum and hence. if the. CRB occurs behind a region of neutral hvdrogen of its host galaxy. provide a direct redshift to à gamma-ray burst.," Assuming the extrapolation of the spectral index $\delta$ ) from the optical to UV passband is correct, we find that a MAMA detector observation of the bright afterglow of a GRB by STIS over $\sim 1$ orbit could reveal a Lyman limit in the spectrum and hence, if the GRB occurs behind a region of neutral hydrogen of its host galaxy, provide a direct redshift to a gamma-ray burst."
+ Given that Lyman limit absorption may come from galaxies along the lineofsight. we predict the expected. frequeney. of Lyman limits in the spectrum as a function of CRB redshift and the distance from their host galaxy.," Given that Lyman limit absorption may come from galaxies along the line–of–sight, we predict the expected frequency of Lyman limits in the spectrum as a function of GRB redshift and the distance from their host galaxy."
+ The authors thank Richard Ale\lahon. Max Pettini and Phillip Outram for insightful comments.," The authors thank Richard McMahon, Max Pettini and Phillip Outram for insightful comments."
+ This paper has also benefitted [rom conversations at the workshop. “Recent Developments Towards Understanding Gamma-Ray Bursts” in Elba. Italy.," This paper has also benefitted from conversations at the workshop, “Recent Developments Towards Understanding Gamma-Ray Bursts” in Elba, Italy."
+ A Lyman limit svstem (LLS) is defined as the region of space for which the optical depth (7) to the Lyman-continuum is ereater than one., A Lyman limit system (LLS) is defined as the region of space for which the optical depth $\tau$ ) to the Lyman-continuum is greater than one.
+ Most Lyman limit svstems are believed to be clouds of LIE with a column density AN(ID>1.671017 env., Most Lyman limit systems are believed to be clouds of HI with a column density $N({\rm HI}) \ge 1.6 \times 10^{17}$ $^{2}$.
+ As depicted in the top half of figure 1. the probability as à function of redshift that a spectrum of a GRB will contain at least one such LLS is where Poisson0LLS[mí(zuiin.2)] is. the Poisson probability of no intervening LLS between the IUD source (at recishift 2) and the observer given the expected number of LLSs: The minimum redshift in which a LES could be detected is cain and NOs) is the number density of LLSs per unit redshift.," As depicted in the top half of figure 1, the probability as a function of redshift that a spectrum of a GRB will contain at least one such LLS is where ${\rm Poisson}\left[0~{\rm LLS} | m(z_{\rm min},z) \right]$ is the Poisson probability of no intervening LLS between the GRB source (at redshift $z$ ) and the observer given the expected number of LLSs: The minimum redshift in which a LLS could be detected is ${z_{\rm
+min}}$ and $N(z)$ is the number density of LLSs per unit redshift."
+ The probability that the redshift) inferred from. an observed. Lyman limit in a GRB afterglow spectrum is at cast c times the redshift of the GRB itself is given as, The probability that the redshift inferred from an observed Lyman limit in a GRB afterglow spectrum is at least $x$ times the redshift of the GRB itself is given as
+The existence of main sequence A- and B-type stars with strong magnetic fields has been known for decades.,The existence of main sequence A- and B-type stars with strong magnetic fields has been known for decades.
+ These peculiar stars (referred to as Ap stars) show anomalous abundances of particular elements. for example of Cr. which may be as much as 10° times (2 dex) overabundant compared to the Sun.," These peculiar stars (referred to as Ap stars) show anomalous abundances of particular elements, for example of Cr, which may be as much as $10^2$ times (2 dex) overabundant compared to the Sun."
+ It is also not uncommon to find atmospheric abundances of rare earths well in excess of solar (?).., It is also not uncommon to find atmospheric abundances of rare earths well in excess of solar \citep{Ryabchikova91}.
+ Almost all Ap stars have angular momenta that are of the order of or less of typical values for normal stars of similar mass., Almost all Ap stars have angular momenta that are of the order of or less of typical values for normal stars of similar mass.
+ The magnetic field strength. spectral line strengths and shapes. and brightness in various photometric bands all vary with the rotation period of the star.," The magnetic field strength, spectral line strengths and shapes, and brightness in various photometric bands all vary with the rotation period of the star."
+ This variability is explained using the oblique rotator model: the magnetic field axis and rotation axis of the Ap star are not aligned with the line of sight. nor with one another. and several elements are distributed non-uniformly over the surface in a pattern that is not axisymmetric about the rotation axis.," This variability is explained using the oblique rotator model: the magnetic field axis and rotation axis of the Ap star are not aligned with the line of sight, nor with one another, and several elements are distributed non-uniformly over the surface in a pattern that is not axisymmetric about the rotation axis."
+ This means that observations through the rotation cycle lead to varying field measurements as a result of observing the magnetic field from different aspects., This means that observations through the rotation cycle lead to varying field measurements as a result of observing the magnetic field from different aspects.
+ Similarly the spectrum. and (as a result of line blocking and backwarming) photometric magnitudes and colours. vary as different parts of the star are observed (?)..," Similarly the spectrum, and (as a result of line blocking and backwarming) photometric magnitudes and colours, vary as different parts of the star are observed \citep{Ryabchikova91}."
+ HD 318107 (=NGC 6405 77). a very peculiar magnetic Ap star. has an effective temperature Των=11800+H-500 K. luminosity logL/L.=1.92+0.1. and mass ΜΜ.0.15 (2).," HD 318107 (=NGC 6405 77), a very peculiar magnetic Ap star, has an effective temperature $\te = 11\,800 \pm 500$ K, luminosity $\log L/L_\odot = 1.92 \pm 0.1$, and mass $M/M_\odot = 2.95 \pm 0.15$ \citep{paper2}."
+. It has a large global magnetic field: the mean magnetic field component lis sometimes as large as 5 or 6 kG. and the typical size of the mean field modulus iis about 15 kG. The star is a member of the open cluster NGC 6405. so that it is known to have an age of logr= (yr).," It has a large global magnetic field; the mean line-of-sight magnetic field component is sometimes as large as 5 or 6 kG, and the typical size of the mean field modulus is about 15 kG. The star is a member of the open cluster NGC 6405, so that it is known to have an age of $\log t = 7.80 \pm 0.15$ (yr)."
+ This means that the star has completed about 17+7% of its main sequence lifetime (??)..," This means that the star has completed about $17 \pm 7$ of its main sequence lifetime \citep{paper2,paper3}."
+ It is still a rather young star., It is still a rather young star.
+ Initially. a rotation period of 52.4 days was proposed by ?: however. more extensive photometric data obtained by ? combined with measurements of the mean surface magnetic field lled to the conclusion that the rotation period is P=9.70850.0021 days.," Initially, a rotation period of 52.4 days was proposed by \citet{North1987}; however, more extensive photometric data obtained by \citet{MM2000} combined with measurements of the mean surface magnetic field led to the conclusion that the rotation period is $P = 9.7085 \pm 0.0021$ days."
+ The magnetic field of this star was modelled by ? based on data from ?.. ? απ further. unpublished magnetic field measurements by Mathys.," The magnetic field of this star was modelled by \citet{LM2000} based on data from \citet{Mathys97}, \citet{MH1997} and further unpublished magnetic field measurements by Mathys."
+ ? adopted a model of colinear magnetic dipole. quadrupole. and octopole components of strength +23700. -23600. and 48300 G respectively.," \citet{LM2000} adopted a model of colinear magnetic dipole, quadrupole, and octopole components of strength +23700, -23600, and +8300 G respectively."
+ Observations of the hemispherically averaged line of sight component of the magnetic field indicate that lis always positive. which leads to the model constraint that the sum of the inclination of the rotation axis to the line of sight. i. and the angle between the magnetic field axis and the rotation axis.B. must be less than about 90 degrees: /-Bx907.," Observations of the hemispherically averaged line of sight component of the magnetic field indicate that is always positive, which leads to the model constraint that the sum of the inclination of the rotation axis to the line of sight, $i$, and the angle between the magnetic field axis and the rotation axis, $\beta$ , must be less than about 90 degrees: $i + \beta \lesssim 90^{\circ}$."
+ This implies that the negative magnetic pole is never directly observed., This implies that the negative magnetic pole is never directly observed.
+ For this model. the values for / and B were chosen to be 11 and 787. with uncertainties discussed in the text of ?..," For this model, the values for $i$ and $\beta$ were chosen to be $11^{\circ}$ and $78^{\circ}$, with uncertainties discussed in the text of \citet{LM2000}."
+ Note that this model does not reproduce the observations very well., Note that this model does not reproduce the observations very well.
+ It appears that the field structure of HD 318107 is more complex than the axisymmetric low order multipole field geometry used. and is probably not even axisymmetric.," It appears that the field structure of HD 318107 is more complex than the axisymmetric low order multipole field geometry used, and is probably not even axisymmetric."
+ Nevertheless. the model provides a first order approximation to the magnetic field geometry of HD318107.," Nevertheless, the model provides a first order approximation to the magnetic field geometry of HD318107."
+ There are rather few hot Ap stars with strong magnetic fields for which detailed chemical abundance analyses have been carried out., There are rather few hot Ap stars with strong magnetic fields for which detailed chemical abundance analyses have been carried out.
+ HD318107. with Zar=11800 K. measured mean field modulus of (B)~15 kG (?) and vsiné=7 km s. is an excellent object for which to earry out such an analysis.," HD318107, with $\te = 11800$ K, measured mean field modulus of $\bs \sim 15$ kG \citep{Mathys97} and $v \sin i \approx 7$ km $^{-1}$ , is an excellent object for which to carry out such an analysis."
+"respectively, leading to an average polarization Pins=0.28+0.03%.","respectively, leading to an average polarization $P_{ins}$ $\pm$."
+. The wavelength range below 3800 is affected by other instrumental problems which make it hardly usable with the typical set of 4 HWP angles (see next section)., The wavelength range below 3800 is affected by other instrumental problems which make it hardly usable with the typical set of 4 HWP angles (see next section).
+" Therefore, this constant correction is sufficient to guarantee the removal of the instrumental polarization with a maximum error of0.05%,, which is comparable to the maximum accuracy one can reach with CAFOS with 4 HWP angles (see next section)."," Therefore, this constant correction is sufficient to guarantee the removal of the instrumental polarization with a maximum error of, which is comparable to the maximum accuracy one can reach with CAFOS with 4 HWP angles (see next section)."
+ We remark that the instrumental polarization correction derived here is strictly valid only for an object placed on the CAFOS reference pixel used for the acquisition onto the 1.0 arcsec slit., We remark that the instrumental polarization correction derived here is strictly valid only for an object placed on the CAFOS reference pixel used for the acquisition onto the 1.0 arcsec slit.
+" With the present analysis we cannot exclude position-dependent effects, similarly to what happens in the FORS instruments (Patat Romaniello 2006))."," With the present analysis we cannot exclude position-dependent effects, similarly to what happens in the FORS instruments (Patat Romaniello \cite{patat06}) )."
+ For the Fourier analysis of the CAFOS polarimetric performances we have used the data obtained for the polarized standard star., For the Fourier analysis of the CAFOS polarimetric performances we have used the data obtained for the polarized standard star.
+ Fig., Fig.
+ 3 shows an example for a 200 wide bin centered at 5500A., \ref{fig:polstd} shows an example for a 200 wide bin centered at 5500.
+". The only components which show a statistically significant power are k=0 and k=4; there is a hint of a nun-null k=2 component, which is related to the so-called pleochroism (Fendt et al. 1996;;"," The only components which show a statistically significant power are $k$ =0 and $k$ =4; there is a hint of a nun-null $k$ =2 component, which is related to the so-called pleochroism (Fendt et al. \cite{fendt};"
+ Patat Romaniello 2006)) but this is only marginally significant at the 5-0 level.," Patat Romaniello \cite{patat06}) ), but this is only marginally significant at the $\sigma$ level."
+" The original signal can be reconstructed using only the k=4 harmonic, with maximum residuals AF; of ~0.1%.."," The original signal can be reconstructed using only the $k$ =4 harmonic, with maximum residuals $\Delta F_i$ of $\sim$."
+ This implies that 4 HWP angles are sufficient to the derive the Stokes parameters with a maximum error of this order., This implies that 4 HWP angles are sufficient to the derive the Stokes parameters with a maximum error of this order.
+ The polarization degree derived using 16 HWP angles at 5500 is 6.4340., The polarization degree derived using 16 HWP angles at 5500 is $\pm$.
+01%.. After applying the instrumental polarization correction described in the previous section this value becomes 6.6+0., After applying the instrumental polarization correction described in the previous section this value becomes $\pm$.
+1%.. This is fully consistent with the reference value 6.70+0.02% measured in the V passband (Schmidt et al. 1992))., This is fully consistent with the reference value $\pm$ measured in the V passband (Schmidt et al. \cite{schmidt}) ).
+ In the example illustrated in Fig., In the example illustrated in Fig.
+" 3 we find ag=1.88+0.01% (the corresponding value derived from the unpolarized standard is 1.9340.01%;; see also Fig. 1,,"," \ref{fig:polstd} we find $a_0$ $\pm$ (the corresponding value derived from the unpolarized standard is $\pm$; see also Fig. \ref{fig:unpol},"
+ upper panel)., upper panel).
+" This indicates that the WP deviates from the ideal case, in that an unpolarized incoming beam is not exactly split into two identical fractions (see Patat Romaniello 2006,, their Sect."," This indicates that the WP deviates from the ideal case, in that an unpolarized incoming beam is not exactly split into two identical fractions (see Patat Romaniello \cite{patat06}, their Sect."
+ 7)., 7).
+" As a consequence, using only 2 HWP angles (which is the minimum set needed to fully reconstruct the Stokes vector) would lead to a very significant error on the final result."," As a consequence, using only 2 HWP angles (which is the minimum set needed to fully reconstruct the Stokes vector) would lead to a very significant error on the final result."
+" To study the instrumental performance as a function of wavelength, we have run the same analysis within 200 wide bins between 3400 and 8600A."," To study the instrumental performance as a function of wavelength, we have run the same analysis within 200 wide bins between 3400 and 8600."
+. The result for the first 6 harmonics is shown in Fig. 4.., The result for the first 6 harmonics is shown in Fig. \ref{fig:fullanal}.
+" The k=0 component is always significant, exceeding ~3% at 7500A,, but this is fairly well corrected if the data set includes at least 4 HWP positions."," The $k$ =0 component is always significant, exceeding $\sim$ at 7500, but this is fairly well corrected if the data set includes at least 4 HWP positions."
+" As for components k=1 and 2, these are detected at a significant level below 3800"," As for components $k$ =1 and 2, these are detected at a significant level below 3800"
+ , 
+"enough. the high-luminosity end of the luminosity function is in better agreement with the pure Y,=0.32 collisions. which is in line with expectations.","enough, the high-luminosity end of the luminosity function is in better agreement with the pure $Y_0 = 0.32$ collisions, which is in line with expectations."
+" The core luminosity function seems to tit slightly better with a pure Y,=0.24 component for luminosities that are one magnitude brighter than the turnoff. which is compatible with mass Segregation since stars of Yo=0.24 are expected to be more massive than stars with ¥y20.32 on average."," The core luminosity function seems to fit slightly better with a pure $Y_0
+= 0.24$ component for luminosities that are one magnitude brighter than the turnoff, which is compatible with mass segregation since stars of $Y_0 = 0.24$ are expected to be more massive than stars with $Y_0
+= 0.32$ on average."
+ The range of colours spanned by our collision models agrees well with the observations. and both place the peak between ACBV)=—0.2 and -0.1.," The range of colours spanned by our collision models agrees well with the observations, and both place the peak between $\Delta(B - V) = -0.2$ and $-0.1$."
+ The observations may fall off more quickly on the blue side. but the difference is within the expected error based on the number of observed stars in the colour bins.," The observations may fall off more quickly on the blue side, but the difference is within the expected error based on the number of observed stars in the colour bins."
+ On the other hand. the observations clearly show an excess of blue stragglers to the red of ACB—V)=-0.I.," On the other hand, the observations clearly show an excess of blue stragglers to the red of $\Delta(B - V) = -0.1$."
+ In our models the stars in this colour range are post-main sequence objects that are in the Hertzsprung gap., In our models the stars in this colour range are post-main sequence objects that are in the Hertzsprung gap.
+ Similar discrepancies between models and observations have been noted beforee., Similar discrepancies between models and observations have been noted before.
+e.mainBodyCitationEnd4577][]SilIs2000. Extra mixing. for instance due to rapid rotation. offers one possible way to extend the lifetime of stars in this region.," Extra mixing, for instance due to rapid rotation, offers one possible way to extend the lifetime of stars in this region."
+ Convective overshooting. as noted before. is ineffective in removing this discrepaney.," Convective overshooting, as noted before, is ineffective in removing this discrepancy."
+ Another possibility is that at least some of the stars in this region are unresolved binaries., Another possibility is that at least some of the stars in this region are unresolved binaries.
+ A K-S test of the cumulative colour distribution gives about equal probability of 0.70 to both model sets B and D. followed by model set A with a probability of 0.62.," A K-S test of the cumulative colour distribution gives about equal probability of $0.70$ to both model sets B and D, followed by model set A with a probability of $0.62$ ."
+" The pure Y,=0.4 models (model set C) has a probability «0.001.", The pure $Y_0 = 0.4$ models (model set C) has a probability $< 0.001$.
+ Combining both the colour and luminosity information. our model set D best describes the blue stragglers in NGC 2808.," Combining both the colour and luminosity information, our model set D best describes the blue stragglers in NGC 2808."
+ None of our luminosity functions fit particularly well., None of our luminosity functions fit particularly well.
+ The best fitting model is model set A. although model set D is not much worse.," The best fitting model is model set A, although model set D is not much worse."
+ By varying the distance modulus. model set A can be made to fit a bit better. while the tit for model set D does not improve significantly.," By varying the distance modulus, model set A can be made to fit a bit better, while the fit for model set D does not improve significantly."
+ The colour distributions for the different model sets are all very similar., The colour distributions for the different model sets are all very similar.
+ The observations show a peak that is 0.15 magnitudes bluer than the turnoff., The observations show a peak that is $0.15$ magnitudes bluer than the turnoff.
+ As with NGC 2808. there are more observed blue stragglers redward of ACB—V)=—O.1 than predicted by our models.," As with NGC 2808, there are more observed blue stragglers redward of $\Delta(B-V) = -0.1$ than predicted by our models."
+ By contrast. no blue stragglers are observed blueward of ΔΙΡV)=—0.3.," By contrast, no blue stragglers are observed blueward of $\Delta(B-V) = -0.3$."
+" The Y,=0.24 models fit slightly better(p= 0.39) than the mixed Y=0.24:0.32 models (p= 0.25) but the picture is not very clear.", The $Y_0 = 0.24$ models fit slightly better $p = 0.39$ ) than the mixed $Y_0=0.24; 0.32$ models $p = 0.25$ ) but the picture is not very clear.
+ Again. none of our models fit very well.," Again, none of our models fit very well."
+ Model set A (Y= 0.24) gives the best agreement with the observations. although model set D is also not a bad fit if we vary the distance modulus.," Model set A $Y_0 = 0.24$ ) gives the best agreement with the observations, although model set D is also not a bad fit if we vary the distance modulus."
+ The theoretical luminosity function falls off perhaps a little too quickly for higher luminosities., The theoretical luminosity function falls off perhaps a little too quickly for higher luminosities.
+ The colour distributions are not very different for the three model sets and all match about equally well., The colour distributions are not very different for the three model sets and all match about equally well.
+ The observations show a peak 0.1 dex bluer than the turnoff that is not present in the models., The observations show a peak 0.1 dex bluer than the turnoff that is not present in the models.
+ The late (post-main sequence) evolution of collision products is interesting. for two reusons., The late (post-main sequence) evolution of collision products is interesting for two reasons.
+ First of all it allows us to test our understanding of the subsequent evolution of collision products by comparing the observed distributions and properties of evolved blue stragglers (post-blue stragglers) with the models., First of all it allows us to test our understanding of the subsequent evolution of collision products by comparing the observed distributions and properties of evolved blue stragglers (post-blue stragglers) with the models.
+ Second of all. it may allow us to probe the dynamical history of the cluster over a longer time interval than can be accomplished with the blue stragglers alone.," Second of all, it may allow us to probe the dynamical history of the cluster over a longer time interval than can be accomplished with the blue stragglers alone."
+ The most promising post-main sequence evolutionary phase to identify stellar collision products is during core helium burning., The most promising post-main sequence evolutionary phase to identify stellar collision products is during core helium burning.
+ Observationally. the collision products are then expected to lie above the horizontal branch(2222).," Observationally, the collision products are then expected to lie above the horizontal branch."
+ We can derive a selection box for evolved collision products by using the evolution models., We can derive a selection box for evolved collision products by using the evolution models.
+" The tirst step is to determine the minimum and maximum values of A4, and B-V during core helium burning for all of our collision models.", The first step is to determine the minimum and maximum values of $M_V$ and $B - V$ during core helium burning for all of our collision models.
+ These define a selection box in the colour-magnitude diagram that contains the core helium burning phase of all our collision products., These define a selection box in the colour-magnitude diagram that contains the core helium burning phase of all our collision products.
+ However. for some models this selection box will now encompass more than just the core helium burning phase.," However, for some models this selection box will now encompass more than just the core helium burning phase."
+ Due to the difference in helium content. he colour of the giant branch is different for different models in our set.," Due to the difference in helium content, the colour of the giant branch is different for different models in our set."
+ We therefore restrict our selection box so that it does not overlap with the red giant branch in any of our models., We therefore restrict our selection box so that it does not overlap with the red giant branch in any of our models.
+ We also impose a cut to remove the cluster horizontal branch., We also impose a cut to remove the cluster horizontal branch.
+ This way. we aim to select only stars that are unambiguously evolved blue stragglers.," This way, we aim to select only stars that are unambiguously evolved blue stragglers."
+ Because we have restricted the selection box. we will not capture all of the core helium burning phase for all collision products. and we may miss ~20% of the core He burning collision products by using this narrower selection box.," Because we have restricted the selection box, we will not capture all of the core helium burning phase for all collision products, and we may miss $\sim 20\%$ of the core He burning collision products by using this narrower selection box."
+ In the logy)Zor / log)L/L; plane our selection box is very similar to that of?., In the $\log_{10} T_\mathrm{eff}$ / $\log_{10} L / L_\odot$ plane our selection box is very similar to that of.
+. We can detine a selection box for the AGB in the same way as for the core helium burning phase. but this selection box turns out to be very narrow.," We can define a selection box for the AGB in the same way as for the core helium burning phase, but this selection box turns out to be very narrow."
+ Since the expected ratio of blue stragglers to AGB post-blue stragglers is also very high. ~1000. it is virtually impossible to compare with observations.," Since the expected ratio of blue stragglers to AGB post-blue stragglers is also very high, $\sim 1000$, it is virtually impossible to compare with observations."
+ Therefore. we only compare the ratio of blue stragglers to core helium burning post-blue stragglers.," Therefore, we only compare the ratio of blue stragglers to core helium burning post-blue stragglers."
+ The predicted ratio of blue stragglers to evolved. blue stragglers is typically 20 — 50., The predicted ratio of blue stragglers to evolved blue stragglers is typically $20$ – $50$.
+ Observationally. the ratios are about 10. but there are only a few stars in the evolved blue straggler selection box and numbers are very sensitive to whether any particular star is classified as an evolved blue straggler or not.," Observationally, the ratios are about $10$, but there are only a few stars in the evolved blue straggler selection box and numbers are very sensitive to whether any particular star is classified as an evolved blue straggler or not."
+ A reasonable agreement between models and observations was reached by based on the average of the horizontal branch and main sequence lifetimes of their collision products. as long as hey selected only the brighter (more massive) collision products.," A reasonable agreement between models and observations was reached by based on the average of the horizontal branch and main sequence lifetimes of their collision products, as long as they selected only the brighter (more massive) collision products."
+ Our approach differs from theirs in two respects: first of all we calculate a population of collision products where each collision oroduet is assigned a weight that depends on the IMF probability of the parent stars as opposed to taking a straight average., Our approach differs from theirs in two respects: first of all we calculate a population of collision products where each collision product is assigned a weight that depends on the IMF probability of the parent stars as opposed to taking a straight average.
+" Second of all we caleulate our population ratios. for both the models and qe observations. based on selection boxes in the colour magnitude ""iagram. not on the relative duration of the evolution phases jemselves."," Second of all we calculate our population ratios, for both the models and the observations, based on selection boxes in the colour magnitude diagram, not on the relative duration of the evolution phases themselves."
+ The first of these has the effect of giving more importance to qe lower mass collision products. which will tend to increase the oediceted ratios.," The first of these has the effect of giving more importance to the lower mass collision products, which will tend to increase the predicted ratios."
+ The effect of using selection boxes that will only capture part of the evolution rather than comparing the lifetimes directly will similarly tend to increase the blue straggler to post-blue straggler ratio., The effect of using selection boxes that will only capture part of the evolution rather than comparing the lifetimes directly will similarly tend to increase the blue straggler to post-blue straggler ratio.
+ Neither of these effects should affect the comparison between our models and the observations. however. because blue stragglers and evolved blue stragglersare selected consistently between the observations and in the models — if he models reflect the observed population. then these population," Neither of these effects should affect the comparison between our models and the observations, however, because blue stragglers and evolved blue stragglersare selected consistently between the observations and in the models – if the models reflect the observed population, then these population"
+Galaxy las a graud spiral striceture. with spira armis. interior structures such as the 3-kpc arms (2?) ancl one or two Galactic bars /3.. ?.. relerences thereiu: ?)).,"Galaxy has a grand spiral structure, with spiral arms, interior structures such as the 3–kpc arms \citep{vanwoerden57, dame08} and one or two Galactic bars \citealt{blitz91}; \citealt{gerhard02}, , references therein; \citealt{benjamin05}) )."
+ Early descriptious of the structure olf our Galaxy were inferred [ro1b 511 «‘outs aud positioune of objects with photometric distauces. but current revisious have beet largely> based upon the kinematics of atorjc and nolecular tracers. stich as and CO.," Early descriptions of the structure of our Galaxy were inferred from star counts and positioning of objects with photometric distances, but current revisions have been largely based upon the kinematics of atomic and molecular tracers, such as and CO."
+" The siuusolcal distribution o. velocities with Galactic ος]udes has established the approximatey circular ane sviumetric rotaion oftl e Galaxy and bohi trace""sl ave been used to coustrain εἶνιamical models (e.g.22???7).."," The sinusoidal distribution of velocities with Galactic longitudes has established the approximately circular and symmetric rotation of the Galaxy and both tracers have been used to constrain dynamical models \citep[e.g.][]{peters75,burton78b,liszt78, binney91,weiner99, fux99}."
+ However. there are limiatiol5 [9] his approach.," However, there are limitations to this approach."
+Hi is detected tuwoughout the Galaxy iu sucl ereal abuidance that only ἰie highest or lowest velocities along a given line of sight (tje terminal velocitjes aud spiral ariital oet1 deviatious) can be used to constaln dyniamical mocleling., is detected throughout the Galaxy in such great abundance that only the highest or lowest velocities along a given line of sight (the terminal velocities and spiral arm tangent deviations) can be used to constrain dynamical modelling.
+ CO is less cliUse and more conceUratecLin the major sirucural features 0 Sour Galaxy. but surveys of CO are lunitecl by a lack of sensitivity to material on the far side o ‘the Galaxy.," CO is less diffuse and more concentrated in the major structural features of our Galaxy, but surveys of CO are limited by a lack of sensitivity to material on the far side of the Galaxy."
+ An alte‘native tracer. the population of 6.7-CGHyz imethanol masers. has the acdvautasesre of only being obseved towards regios of high-mass star formaion (????7) auc the strong aid narrow. emissiou o ‘these masers is well correlated with their syseimic velocity. the central maser veocity typically within 1 (??)..," An alternative tracer, the population of 6.7–GHz methanol masers, has the advantage of only being observed towards regions of high-mass star formation \citep{pestalozzi02b,minier03,sridharan02,walsh03,xu08} and the strong and narrow emission of these masers is well correlated with their systemic velocity, the central maser velocity typically within $^{-1}$ \citep{szymczak07, pandian09}."
+ They are also found in abundance throughout the Galaxy {/2) aud have been detected to Calactocentrie distances of kkpe (/2).., They are also found in abundance throughout the Galaxy \citep{pestalozzi05} and have been detected to Galactocentric distances of kpc \citep{honma07}.
+ Inuer Galaxy structires |ave been shown to be traced by 6.77GHz inethauol masers: 15 sources are associated witty tje 2—MS-kpe arms (4(2 )z T sources with the hieh velocities associated with Calactic bars (222): and Ll « ahsagittarius B2 (?.relereucesthereiu)..," Inner Galaxy structures have been shown to be traced by 6.7–GHz methanol masers: 45 sources are associated with the 3–kpc arms \citep{green10mmb2}; ; 7 sources with the high velocities associated with Galactic bars \citep{caswell97, caswell10mmb1,green10mmb2}; and 11 with Sagittarius B2 \citep[][references therein]{caswell10mmb1}."
+ The 6.7-CHz metliauol masers are excellent measures o ‘the slnematic behaviour of the structure of our Galaxy within the terminal velocities outlined in aiid CO., The 6.7–GHz methanol masers are excellent measures of the kinematic behaviour of the structure of our Galaxy within the terminal velocities outlined in and CO.
+ Iu this paper we explo‘e how 6.7-GHz inehanol iuasers trace the uajor sta‘forming structures ol the inner Galaxy. specifically: the Galactic |xw. the 3 kpe aruis. aud the origins of the spiral arms.," In this paper we explore how 6.7–GHz methanol masers trace the major star forming structures of the inner Galaxy, specifically: the Galactic bar, the 3 kpc arms, and the origins of the spiral arms."
+ All of these structures are believed o be confined within 1 kpe of tje Galacti€ centre., All of these structures are believed to be confined within $\sim$ 4 kpc of the Galactic centre.
+ Assuming a solar distauce of 8.1 kpe (?).. the inner | spe of the Galaxy is geometricaly contained. witlin he Galactic longitude rauge —28* o +28°.," Assuming a solar distance of 8.4 kpc \citep{ghez08}, the inner 4 kpc of the Galaxy is geometrically contained within the Galactic longitude range $-$ $^{\circ}$ to $+$ $^{\circ}$."
+ The Methanol Multibeau (MMB survey (2).. which observed from —1T1 to +60° wit1 the Parkes Raclio Telescope. «elected. 7550 masers within ougitudes E287.," The Methanol Multibeam (MMB) survey \citep{green09a}, which observed from $-$ $^{\circ}$ to $+$ $^{\circ}$ with the Parkes Radio Telescope, detected $\sim$ 550 masers within longitudes $\pm$ $^{\circ}$."
+ The MMB had a 3 σ seusitivity of 0.7 Jy. covered tle latituke rauge of 42° and rad a velocity coverage euconmpassii the extent of CO emission seen in ?..," The MMB had a 3 $\sigma$ sensitivity of 0.7 Jy, covered the latitude range of $\pm$ $^{\circ}$ and had a velocity coverage encompassing the extent of CO emission seen in \citet{dame01}."
+ The MAB catalogue herefore reoreseuts the most compele survey of current high-mass 5ar fornuion traced by methanol masers in the region. providiug a ceusus of the regions wich inelides their kinematic jhaviour.," The MMB catalogue therefore represents the most complete survey of current high-mass star formation traced by 6.7-GHz methanol masers in the region, providing a census of the regions which includes their kinematic behaviour."
+ Through examining the distribution aud densities of these masers in longitude aud velocity. in comparisou with the spiral arms. we attempt to clelineate he inner Calactic structure.," Through examining the distribution and densities of these masers in longitude and velocity, in comparison with the spiral arms, we attempt to delineate the inner Galactic structure."
+ The longitude-velocity(Le) distribution of 6.7-GHz methanol masers betweenlougitucles 28* is shown in, The longitude-velocity) distribution of 6.7–GHz methanol masers betweenlongitudes $\pm$ $^{\circ}$ is shown in
+where the negative solutions are excluded.,where the negative solutions are excluded.
+ Similarly when 0=7/2. that in the latter case the MIID wave couldn't propagate very far in the accretion dise. so we consider the coupling modes that propagate along the z axis (0=0) as a more promising explanation of the QPOs.," Similarly when $\theta
+=\pi/2$, Note that in the latter case the MHD wave couldn't propagate very far in the accretion disc, so we consider the coupling modes that propagate along the $z$ axis $\theta=0$ ) as a more promising explanation of the QPOs."
+ We first rule out the possibility of the ws mode as the lower kIIz QPOs. since Eq. (," We first rule out the possibility of the $\omega_2$ mode as the lower kHz QPOs, since Eq. ("
+21) requires (hat it should always higher than © for stable accretion. which is contracdicted with observations.,"21) requires that it should always higher than $\Omega$ for stable accretion, which is contradicted with observations."
+ The coupling between à and i; can also be excluded. which implies a constant [requency separation.," The coupling between $\omega_1$ and $\omega_3$ can also be excluded, which implies a constant frequency separation."
+ So we suggest (he upper ancl the lower kIIz QPO [requencies be ve=waf/252 and vy=ποπ.," So we suggest the upper and the lower kHz QPO frequencies be $\nu _2 =
+\omega_2 / 2\pi$ and $\nu_1 = \omega_1 / 2\pi$."
+ From Eqs. (, From Eqs. (
+"20) and (21) we can get the following relation between the Ireequencies of (he upper and the lower kIIz QPOs. where v,=O/27. or where 7?=(ip—A7)/(1+A7) and 0?=(AP77)/(1477).","20) and (21) we can get the following relation between the frequencies of the upper and the lower kHz QPOs, where $\nu_s= \Omega / 2\pi$, or where $\varepsilon^2 = (\eta ^2 - \lambda ^2)/(1 + \lambda ^2)$ and $\delta^2 = (\lambda ^2 - \eta ^2)/(1 + \eta ^2)$."
+" Equations (27) and (28) indicate that the (win kIIlz QPOs may be divided into two groups. with the slope of the 15/1, vs. vy/vs relation either larger or smaller (han 1."," Equations (27) and (28) indicate that the twin kHz QPOs may be divided into two groups, with the slope of the $\nu_2/\nu_s$ vs. $\nu_1/\nu_s$ relation either larger or smaller than 1."
+ In the following we call them the laree slope coellicient sources (LSCS) and (he small slope coefficient sources (SSCS). respectively.," In the following we call them the large slope coefficient sources (LSCS) and the small slope coefficient sources (SSCS), respectively."
+ We compare in Fig., We compare in Fig.
+" 2 the 1s/r, vs. i/i, relations obtained in last section with (he observed kIIz QPOs in six sources 4U 06144-09. 4U 1608—252. 4U 1636-53. 4U 1728-34. 4U 1915—05."," 2 the $\nu_2/\nu_s$ vs. $\nu_1/\nu_s$ relations obtained in last section with the observed kHz QPOs in six sources 4U $+$ 09, 4U $-$ 52, 4U $-$ 53, 4U $-$ 34, 4U $-$ 05,"
+ancl is boosted by Ον=E.,and is boosted by ${\cal D} (t_p)= \Gamma$.
+" Alter /,. emission arrives from 0p +. for which /ME(1787|1)/2 and on virtue of equations (6)) and (7))."," After $t_p$, emission arrives from $\theta > 
+\Gamma^{-1}$ , for which $t/t_p = (\Gamma^2\theta^2 + 1)/2$ and on virtue of equations \ref{rt}) ) and \ref{D}) )."
+" '""herefore. the large-angle emission. arriving at. fixed observer [requeney corresponds to. an ever-increasing comoving-Erame frequency."," Therefore, the large-angle emission arriving at fixed observer frequency corresponds to an ever-increasing comoving-frame frequency."
+ Then. the above conclusion that. [or GR 990123. optical is below the peak-energy. of the synchrotron spectrum (ον) implies that the svynchrotron light-curve should. steepen. at { when ορ crosses. the optical domain: where prime denotes a quantity in the comovine-frame.," Then, the above conclusion that, for GRB 990123, optical is below the peak-energy of the synchrotron spectrum $\epsilon_{p,sy}$ ) implies that the synchrotron light-curve should steepen at $t_+$ when $\epsilon_{p,sy}$ crosses the optical domain: where prime denotes a quantity in the comoving-frame."
+ The ROTSE emission with the forward-shock contribution subtracted (Figure 2)) shows such a steepening at about 45 s. the subsequent decay index. ax2. implying an optical. SED slope 320. as expected at the peak of synchrotron spectrum.," The ROTSE emission with the forward-shock contribution subtracted (Figure \ref{t30}) ) shows such a steepening at about 45 s, the subsequent decay index, $\alpha \simeq 2$, implying an optical SED slope $\beta_o \simeq 0$, as expected at the peak of synchrotron spectrum."
+ For now. we parameterize the time when/ relative to the epoch of the second GRB pulse peak: [oΞαν.," For now, we parameterize the time when $t_+$ relative to the epoch of the second GRB pulse peak: $t_+ = x\, t_p$."
+ TPhorefore. the observer-Lrame svnchrotron peak-enereyv is The extrapolation of the power-law fit to the ROSE light-curve to ἐν. predicts an optical Dus of fn(4)=4 Jv.," Therefore, the observer-frame synchrotron peak-energy is The extrapolation of the power-law fit to the ROTSE light-curve to $t_p$, predicts an optical flux of $F_R (t_p) = 4$ Jy."
+" The extrapolation of the forward-shock (ES) emission to the same time is resi.)=0.98 Js. therefore the peak of the excess optical emission. (which we attribute to the same mechanism as the burstitself) is Εν)=Fn)Brstt,)=3.07 Jv (the forward-shock emission may have started later than the optical peak time. in which case the εν) above underestimates the true svnchrotron [ux by ))."," The extrapolation of the forward-shock (FS) emission to the same time is $F_{FS} (t_p) = 0.93$ Jy, therefore the peak of the excess optical emission (which we attribute to the same mechanism as the burstitself) is $F_{sy} (t_p) 
+= F_R(t_p)-F_{FS}(t_p) = 3.07$ Jy (the forward-shock emission may have started later than the optical peak time, in which case the $F_{sy} (t_p)$ above underestimates the true synchrotron flux by )."
+" Taking into account that the slope of the optical SED is 3,=1/3. it follows that the Hux at the peak of the svnchrotron spectrum is The Hux Baa at the peak-energy ον Κον of the inverse-C'ompton spectrum can be derived from the 100 keV Hux of 0.61 mJv reported by Briges (1999) at 17 s alter the onset ofthe second GRB pulse and from the burst Dux decrease by about from f,= Ss to l7s: Figure 3. shows schematically the svachrotron and inverse-C'ompton. SED. the characteristics of the former having been derived above in the larec-angle emission interpretation of the optical counterpart. while those of the latter come clirectly from observations."," Taking into account that the slope of the optical SED is $\beta_o = 1/3$, it follows that the flux at the peak of the synchrotron spectrum is The flux $F_{p,ic}$ at the peak-energy $\epsilon_{p,ic}$ keV of the inverse-Compton spectrum can be derived from the 100 keV flux of 0.61 mJy reported by Briggs (1999) at 17 s after the onset of the second GRB pulse and from the burst flux decrease by about from $t_p=8$ s to 17 s: Figure \ref{spek} shows schematically the synchrotron and inverse-Compton SED, the characteristics of the former having been derived above in the large-angle emission interpretation of the optical counterpart, while those of the latter come directly from observations."
+ Equations (3)). (52). (100). CL1)) and (12)) represent the conditions which we useto constrain the physical parameters of the svynchrotron. self-Compton. model for the IOTSIE optical and BATSE 5-rav. emissions of GRB 990123.," Equations \ref{epic}) ), \ref{gc}) ), \ref{epsy}) ), \ref{fpsy}) ) and \ref{fpic}) ) represent the conditions which we use to constrain the physical parameters of the synchrotron self-Compton model for the ROTSE optical and BATSE $\gamma$ -ray emissions of GRB 990123."
+" An upper limit on the parameter ac is obtained if the svnchrotron power-law spectrum above εν extends up to the 5-ray range. there are no other spectral breaks but εκ. by requiring that the svnchrotron flux at 2 keV. (the lowest X-ray observational [requenev) does not exceed the inverse-Compton Lux: which leads tour<14 for cosy—Tei. which is the lowest value of the cooling energy. (see equation 5)). and to roxὉ for ou©2 keV. Therefore. under the assumption that the svncehrotron spectrum extends up keV energies. the peak-energv of the svnchrotron spectrum should. satisfy and should cross the optical at /—r£,«115 s after the peak of the second. CRB pulse (145 s after trigeer)."," An upper limit on the parameter $x$ is obtained if the synchrotron power-law spectrum above $\epsilon_{p,sy}$ extends up to the $\gamma$ -ray range, there are no other spectral breaks but $\epsilon_{c,sy}$, by requiring that the synchrotron flux at 2 keV (the lowest $X$ -ray observational frequency) does not exceed the inverse-Compton flux: which leads to $x < 14$ for $\epsilon_{c,sy} = 7\epsilon_{p,sy}$, which is the lowest value of the cooling energy (see equation \ref{gc}) ), and to $x < 9$ for $\epsilon_{c,sy} 
+> 2$ keV. Therefore, under the assumption that the synchrotron spectrum extends up keV energies, the peak-energy of the synchrotron spectrum should satisfy and should cross the optical at $t_+ = x\,t_p < 115$ s after the peak of the second GRB pulse (145 s after trigger)."
+ Lf ROTVSE emission is mostly the large-angle emission of the second GRB pulsethen the optical lightcurve should exhibit a break at //. which is consistent with observations (see Figure 2)).," If ROTSE emission is mostly the large-angle emission of the second GRB pulsethen the optical light-curve should exhibit a break at $t_+$ , which is consistent with observations (see Figure \ref{t30}) )."
+ Based on the constraint above. we normalize .r to 10.," Based on the constraint above, we normalize $x$ to 10."
+ Integrating the scattering photon spectrum. per. electron equation 2.48 of Blumenthal Could. 1970) over the, Integrating the scattering photon spectrum per electron equation 2.48 of Blumenthal Gould 1970) over the
+for auv we.,for any $\omega$.
+ We use for the massless scalar field the generalforum o(7..0)=ar|For). where ¢=coustaut.," We use for the massless scalar field the generalform $\phi(\tau,x)=a\tau+F(x)$ , where $a$ =constant."
+ To solve the above equations in the general case is not an easy task., To solve the above equations in the general case is not an easy task.
+ In the following. we shall consider several particular cases.," In the following, we shall consider several particular cases."
+ We first show some particular solutious for a dust fluid aud later for thegeneral case With οz0., We first show some particular solutions for a dust fluid and later for thegeneral case with $\omega\neq0$ .
+ Solutions corresponding to p=—0. that have already been considered in reference [23].. are also shown in section 3.3.," Solutions corresponding to $p=\rho=0$, that have already been considered in reference \cite{HWW04}, are also shown in section 3.3."
+ From geodesie equatious we find that (ce)=&y. constant.," From geodesic equations we find that $\Phi(x)=\Phi_0$, constant."
+" Then. the Eiusteiu equations are given by Lo.) Sto, py boy) 10,0] το woJJ20."," Then, the Einstein equations are given by ] ] _0 ] ] ] ] = 0."
+0 Iu the following. we consider several particular cases.," In the following, we consider several particular cases."
+ where $9. Vy. Syaud oy are arbitrary integration constants.," where $\Phi_0$ , $\Psi_0$, $S_0$and $\phi_0$ are arbitrary integration constants."
+also about 15.,also about 1.
+.. All our variables used. in these tests are independent of each other. thus the probability to get all five xvwameters fitted once by the random supercluster mocdoel simultaneously is much smaller than 1%.," All our variables used in these tests are independent of each other, thus the probability to get all five parameters fitted once by the random supercluster model simultaneously is much smaller than 1."
+. Even if using the random supercluster model is a fast rut not ideal procedure for calculating these probabilities he main result would be hardly changed by more ingenious simulations: the probability is very small., Even if using the random supercluster model is a fast but not ideal procedure for calculating these probabilities the main result would be hardly changed by more ingenious simulations: the probability is very small.
+ hus we conclude hat within standard cosmological mocels it is dillicult to eencrate the observed correlation function., Thus we conclude that within standard cosmological models it is difficult to generate the observed correlation function.
+ Which perturbation spectrum can produce the observed correlation function of clusters in. rich superclusters?, Which perturbation spectrum can produce the observed correlation function of clusters in rich superclusters?
+ Analytic calculations. made in Paper HE show that the correlation function has an oscillatory behaviour only if he power spectrum has a peak at the corresponding waventunber., Analytic calculations made in Paper III show that the correlation function has an oscillatory behaviour only if the power spectrum has a peak at the corresponding wavenumber.
+ In this paper it was also demonstrated that the sharpness and height of the peak in the spectrum determines he character of oscillations of the correlation function., In this paper it was also demonstrated that the sharpness and height of the peak in the spectrum determines the character of oscillations of the correlation function.
+ Llere we estimate the possible shape of the spectrum. on scales of interest using comparison with models with known spectra., Here we estimate the possible shape of the spectrum on scales of interest using comparison with models with known spectra.
+" We shall compare the spectra and correlation ""unctions of three models: the standard CDM model. the double. power-law model. and a mixed. geometrical. moclel consisting of two populations. one with superclusters located randomly along regularlv spaced. rods. anc the other of irregularly spaced. superclusters (see Paper HE for details)."," We shall compare the spectra and correlation functions of three models: the standard CDM model, the double power-law model, and a mixed geometrical model consisting of two populations, one with superclusters located randomly along regularly spaced rods and the other of irregularly spaced superclusters (see Paper III for details)."
+ Power spectra of these three models are shown in Figure 9., Power spectra of these three models are shown in Figure 9.
+ We see that the double-power law model and the mixed mocel have rather similar spectra near the maximum., We see that the double-power law model and the mixed model have rather similar spectra near the maximum.
+ Both models have also similar correlation functions with weak oscillations (see Ligure 4 of Paper UL)., Both models have also similar correlation functions with weak oscillations (see Figure 4 of Paper III).
+ Phe oscillations are more regular in the ecometrical model. as expected.," The oscillations are more regular in the geometrical model, as expected."
+ However. the dillerences between models are not large.," However, the differences between models are not large."
+ The maximal deviation of the spectrum near the maximum from he corresponding CDAM-tvpe spectrum. is 0.2 dex. Le. about a factor of 1.25 in amplitudo.," The maximal deviation of the spectrum near the maximum from the corresponding CDM-type spectrum is $\sim 0.2$ dex, i.e. about a factor of 1.25 in amplitude."
+" ""These models show that already ài mocest: deviation rom the standard. CDM spectrum. produces an oscillating correlation function for clusters in rich superclusters.", These models show that already a modest deviation from the standard CDM spectrum produces an oscillating correlation function for clusters in rich superclusters.
+ The actual power spectrum of our cluster sample has a »ak of even higher amplitude (see LOT)., The actual power spectrum of our cluster sample has a peak of even higher amplitude (see E97).
+ We have determined the correlation function for clusters of ealaxies separately for all clusters and for clusters located in rich and in poor superclusters., We have determined the correlation function for clusters of galaxies separately for all clusters and for clusters located in rich and in poor superclusters.
+ The correlation function of clusters in rich superclusters that form the skeleton of the supercluster-void network has an oscillatory behaviour with a period of 145+15Mpe..," The correlation function of clusters in rich superclusters that form the skeleton of the supercluster-void network has an oscillatory behaviour with a period of $115
+\pm 15$."
+ Within an interval of ~650 oover which observational data are available. five secondary maxima and minima of the correlation function are seen.," Within an interval of $\sim650$ over which observational data are available, five secondary maxima and minima of the correlation function are seen."
+ The amplitude of oscillations is larger lor clusters located in very rich superclusters., The amplitude of oscillations is larger for clusters located in very rich superclusters.
+ The scale of the supercluster-void network found. here on the basis of the cluster correlation function is rather close to the scale found using other methods. such as void cliameter analysis. pencil-beam studies. or absorbers in the inc-of-sight to QSOs (Quashnock et al.," The scale of the supercluster-void network found here on the basis of the cluster correlation function is rather close to the scale found using other methods, such as void diameter analysis, pencil-beam studies, or absorbers in the line-of-sight to QSOs (Quashnock et al."
+ 1996). although the alter applies to higher redshifts.," 1996), although the latter applies to higher redshifts."
+ The reality of oscillations of the cluster correlation unction is supported by the following arguments. (, The reality of oscillations of the cluster correlation function is supported by the following arguments. (
+1) The error corridor of the correlation function etermined For clusters in rich superclusters is much smaller jin the amplitude of oscillations. (,1) The error corridor of the correlation function determined for clusters in rich superclusters is much smaller than the amplitude of oscillations. (
+2) Oscillations are seen in cluster samples located in both Galactic hemispheres. (,2) Oscillations are seen in cluster samples located in both Galactic hemispheres. (
+3) Similar oscillations with lower amplitude are observed. in 16 Las Campanas Redshift Survey of galaxies by Tucker (1995. 1997). (,"3) Similar oscillations with lower amplitude are observed in the Las Campanas Redshift Survey of galaxies by Tucker (1995, 1997). ("
+4) In all samples the shape of the oscillating correlation function follows almost exactly the expected shape for a quas-regular network of superclusters ancl voids. (,4) In all samples the shape of the oscillating correlation function follows almost exactly the expected shape for a quasi-regular network of superclusters and voids. (
+5) The double conical shape of the volume sampled. by clusters cannot influence the results. (,5) The double conical shape of the volume sampled by clusters cannot influence the results. (
+6) Parameters of the oscillations practically co not depend. on the smoothing length of the correlation function. nor on the neighbourhood radius used in supercluster definition. or on errors of the selection function used to calculate the correlation function.,"6) Parameters of the oscillations practically do not depend on the smoothing length of the correlation function, nor on the neighbourhood radius used in supercluster definition, or on errors of the selection function used to calculate the correlation function."
+ The correlation length of clusters of galaxies depends on the cluster population: for clusters in poor superclusters it is about 17Alpe:: for clusters in rich superclusters it is about 45Mpc., The correlation length of clusters of galaxies depends on the cluster population: for clusters in poor superclusters it is about 17; for clusters in rich superclusters it is about 45.
+. We have compared. the observed. correlation. function with correlation functions calculated for. clusters in CDALmodcels anc for models with randomly: clistributect superclusters., We have compared the observed correlation function with correlation functions calculated for clusters in CDM-models and for models with randomly distributed superclusters.
+ These. models have a broad-band power spectrum with a smooth transition between the positive spectral index at long wavelengths and a negative index ab small wavelengths., These models have a broad-band power spectrum with a smooth transition between the positive spectral index at long wavelengths and a negative index at small wavelengths.
+ In these models. the correlation function of clusters in rich. superclusters located. in double, In these models the correlation function of clusters in rich superclusters located in double
+two outer lobes of a radio source. and the nearest component is a feature in the jet of one of the sources]: tiii) the nearest component is within 15 aresec of the optical galaxy. with the second and third components both within 90 aresec and the angle NVSS2-NVSSI-NVSS3 greater than 135 degrees [this is the case where the nearest component corresponds to the core of the triple: the offset and angle classification requirements distinguish this from a single component source with two unassociated sources].,"two outer lobes of a radio source, and the nearest component is a feature in the jet of one of the sources]; (iii) the nearest component is within 15 arcsec of the optical galaxy, with the second and third components both within 90 arcsec and the angle NVSS2-NVSS1-NVSS3 greater than 135 degrees [this is the case where the nearest component corresponds to the core of the triple: the offset and angle classification requirements distinguish this from a single component source with two unassociated sources]."
+ Galaxies which satistied these constraints were investigated using FIRST to accept or reject obvious cases., Galaxies which satisfied these constraints were investigated using FIRST to accept or reject obvious cases.
+ Galaxies were accepted as triples if they possessed a FIRST source within 3 arcsec of the optical galaxy position., Galaxies were accepted as triples if they possessed a FIRST source within 3 arcsec of the optical galaxy position.
+ They were rejected if. as for the doubles. they had 3 or fewer FIRST components. all further than 15 aresee from the optical galaxy. with total flux equal to at least half of the sum of the three NVSS fluxes.," They were rejected if, as for the doubles, they had 3 or fewer FIRST components, all further than 15 arcsec from the optical galaxy, with total flux equal to at least half of the sum of the three NVSS fluxes."
+ The remainder of the galaxies were referred to visual analysis., The remainder of the galaxies were referred to visual analysis.
+ Figure 2 compares the results of this analysis for the SDSS galaxies and the random sample., Figure \ref{nvsstrpls} compares the results of this analysis for the SDSS galaxies and the random sample.
+ Galaxies that were neither classified as triples nor visually inspected were then investigated to see if they were associated with a double radio source., Galaxies that were neither classified as triples nor visually inspected were then investigated to see if they were associated with a double radio source.
+ Each of the three NVSS pairs was checked using the double source analysis described above., Each of the three NVSS pairs was checked using the double source analysis described above.
+ Galaxies with + or more NVSS sources within 3 areminutes of the galaxy are likely either to be the host galaxy of a multiple-component radio source. or to lie close to one.," Galaxies with 4 or more NVSS sources within 3 arcminutes of the galaxy are likely either to be the host galaxy of a multiple-component radio source, or to lie close to one."
+ No such cases were accepted without visual analysis., No such cases were accepted without visual analysis.
+ Visual analysis was carried out on all galaxies with 5 or more NVSS matches. as well as on galaxies with 4 NVSS matches where either the mean position or the flux-weighted mean position of the + NVSS sources was within 30 aresee of the optical galaxy.," Visual analysis was carried out on all galaxies with 5 or more NVSS matches, as well as on galaxies with 4 NVSS matches where either the mean position or the flux-weighted mean position of the 4 NVSS sources was within 30 arcsec of the optical galaxy."
+ All other galaxies were not considered to be quadrupoles. but," All other galaxies were not considered to be quadrupoles, but"
+allowed in and numerical models (Tchekhovskoy 2009).,"allowed in analytical and numerical models \citep[]{T08, L09, K09}."
+". X-ray observations reveal the ISM properties such as the Bondi radius rg~ 250 pc and the King core radius Το~1.4 kpc (see,e.g.,,2006). Youn"," X-ray observations reveal the ISM properties such as the Bondi radius $r_{\rm B} \sim$ 250 pc and the King core radius $r_{\rm c} \simeq 1.4$ kpc \citep[see, {\em e.g.},."
+"gThus, the region of conical streamlines in the M87 jet lies between rp and the marginal radius for the power-law decay beyond r, in the King profile, indicating that the ISM distribution is essentially uniform."," Thus, the region of conical streamlines in the M87 jet lies between $r_{\rm B}$ and the marginal radius for the power-law decay beyond $r_{\rm c}$ in the King profile, indicating that the ISM distribution is essentially uniform."
+ We thus rule out that the structure downstream of HST-1 is hydrostatically confined by Pism in order to conform to a conical streamline.," We thus rule out that the structure downstream of HST-1 is hydrostatically confined by $p_{\rm
+ism}$ in order to conform to a conical streamline."
+" In order to possess a conical streamline without any collimation between knots HST-1 and A, the condition pjetDism Should be maintained."," In order to possess a conical streamline without any over-collimation between knots HST-1 and A, the condition $ p_{\rm jet} \gtrsim p_{\rm
+ism}$ should be maintained."
+ Knots HST-1 to A do appear to be over-pressured with respect to the external pressure (Owenetal. 1989)., Knots HST-1 to A do appear to be over-pressured with respect to the external pressure \citep[]{O89}.
+". However, of the inter-knot as estimated by the minimum Ῥιοιenergy argument for the regions,VLA data (Sparksetal.1996),, appears under-pressured with respect to Pism, as estimated by the recent X-ray observations (Youngetal.2002;Rafferty 2006).."," However, $p_{\rm jet}$ of the inter-knot regions, as estimated by the minimum energy argument for the VLA data \citep[]{S96}, appears under-pressured with respect to $p_{\rm ism}$, as estimated by the recent X-ray observations \citep[]{YWM02, R06}. ."
+" One possibility is an underestimation of the magnetic field strength when the toroidal (azimuthal) components are not considered, cf. Owenetal.(1989)."," One possibility is an underestimation of the magnetic field strength when the toroidal (azimuthal) components are not considered, cf. \citet[]{O89}."
+. It has been further suggested that the magnetic field energy is at least in equipartition (or even larger with a factor of 1 ~ 2) with the energy of the radiating ultrarelativistic electrons (Stawarzetal.2005)., It has been further suggested that the magnetic field energy is at least in equipartition (or even larger with a factor of 1 $\sim$ 2) with the energy of the radiating ultrarelativistic electrons \citep[]{S05}.
+. We note that over-pressured knots do appear to have trails of stationary recollimation shocks in purely hydrodynamic jets., We note that over-pressured knots do appear to have trails of stationary recollimation shocks in purely hydrodynamic jets.
+ Falle&Wilson(1985) performed hydrodynamic simulations to apply stationary recollimation shocks to the observed knots at VLA scales under the assumption of the shallow ISM gradient (piacx z-l)," \citet[]{FW85}
+ performed hydrodynamic simulations to apply stationary recollimation shocks to the observed knots at VLA scales under the assumption of the shallow ISM gradient $p_{\rm ism}\propto z^{-1}$ )."
+" Stationary features, however, are in conflict with the observed large proper motions (Biretta1995, 1999),, while the ISM also does not appear to have such a gradient."," Stationary features, however, are in conflict with the observed large proper motions \citep[]{B95, B99}, while the ISM also does not appear to have such a gradient."
+" Therefore, we suggest that the highly magnetized nature of the jet may be responsible for the conical part of the M87 jet."," Therefore, we suggest that the highly magnetized nature of the jet may be responsible for the conical part of the M87 jet."
+ We next consider a parabolic streamline (1«a< 2) for supersonic jets., We next consider a parabolic streamline $1 < a \le 2$ ) for supersonic jets.
+ It is shown thatthe magnetized jet can be parabolic in analytical and numerical models where the ISM, It is shown thatthe magnetized jet can be parabolic in analytical and numerical models where the ISM
+orders of magnitude smaller Caan (hose associated with the effects of partial ionisation so they have been neglected in this paper.,orders of magnitude smaller than those associated with the effects of partial ionisation so they have been neglected in this paper.
+ We assume (he medium to be an ideal gas., We assume the medium to be an ideal gas.
+" The equations are non-dimentionalized using the sound speed (C;=13.2 1) ). the pressure scale height (A=C.ίσα)BR,T/(ug)6.1x10* em). the density at the centre of the prominence (pr—0)210P gem *) giving a characteristic timescale of 7=A/C.41 ss. The value lor >=1.05 is taken. (his is for simplicity as the internal energy e=p/(5—1) excludes the possibility of using ?=1."," The equations are non-dimentionalized using the sound speed $C_s=13.2$ $^{-1}$ ), the pressure scale height $\Lambda=C_s/(\gamma g)= R_gT/(\mu g)=6.1 \times 10^7$ cm), the density at the centre of the prominence $\rho (x=0)=10^{-13}$ g $^{-3})$ giving a characteristic timescale of $\tau=\Lambda/C_s=47$ s. The value for $\gamma=1.05$ is taken, this is for simplicity as the internal energy $\epsilon=p/(\gamma-1)$ excludes the possibility of using $\gamma=1$ ."
+ Using 9=0.5. which gives an Allvénn velocity of Vy=CLὃς b ," Using $\beta=0.5$, which gives an Alfvénn velocity of $V_A=C_s\sqrt{2/\gamma \beta}=25.8$ $^{-1}$."
+"The initial model is as follows: where 2,9 is the value of the horizontal field al #=0 and Ό.ς is the value of the verlical field) as i—oc.", The initial model is as follows: where $B_{x0}$ is the value of the horizontal field at $x=0$ and $B_{z\infty}$ is the value of the vertical field as $x \rightarrow \infty $.
+ This model is (he Kippenhahn-5chlütter model as presented in Priest(1982)., This model is the Kippenhahn-Schlütter model as presented in \cite{PR1982}.
+. As the Nippenhahn-Schlütter model is linearly stable to ideal ΑΠΟ perturbations. a nonlinear perturbation is necessary.," As the Kippenhahn-Schlütter model is linearly stable to ideal MHD perturbations, a nonlinear perturbation is necessary."
+ The perturbation considered here is a high temperature. low densitv tube placed in the center of the Ixippenhahn-»chlütter model.," The perturbation considered here is a high temperature, low density tube placed in the center of the Kippenhahn-Schlütter model."
+ The density of the bubble at 7=z0 is 0.3p(0) (Atwood number 4= 0.53) with width 2X and height SA., The density of the bubble at $x=z=0$ is $0.3\rho(0)$ (Atwood number $A=0.53$ ) with width $2\Lambda$ and height $8\Lambda$.
+ Figure 2. shows a visual representation of the iniGal conditions., Figure \ref{init} shows a visual representation of the initial conditions.
+ The color contour represents (he mass densitv. the lines represent (he magnetic field lines.," The color contour represents the mass density, the lines represent the magnetic field lines."
+ The initially the physical quantities are uniform in the y direction., The initially the physical quantities are uniform in the $y$ direction.
+" To excite the interchange mode a velocity perturbation in e, was imposed. where v, was given as a sunm of sinusoidal curves of dillerent wavelength."," To excite the interchange mode a velocity perturbation in $v_y$ was imposed, where $v_y$ was given as a sum of sinusoidal curves of different wavelength."
+ The maximum amplitude of the perturbation (|t'u ) is less than 0.01C*.., The maximum amplitude of the perturbation $|v_y|$ ) is less than $0.01 C_s$.
+ To reduce computational time. we assume a rellective symmetry boundary. at 6=0.," To reduce computational time, we assume a reflective symmetry boundary at $x=0$."
+" Due to the nature of the magnetic field at the top and bottom (2) boundary aud at à=L,. the choice for boundary is verv Hlimited."," Due to the nature of the magnetic field at the top and bottom $z$ ) boundary and at $x=L_x$, the choice for boundary is very limited."
+" A free boundary. is assumed al 7=L, with a damping zone (damping time 7= 4.4) for the hydromagnetic variables aud D. (to maintain the angle of the magnetic field al the boundary).", A free boundary is assumed at $x=L_x$ with a damping zone (damping time $\tau=4.4$ ) for the hydromagnetic variables and $B_z$ (to maintain the angle of the magnetic field at the boundary).
+ This condition is important to maintain tension at the boundary (to make sure the plasma can besupported there. but as (his is a local simulation this will ultimately have an influence the results.," This condition is important to maintain tension at the boundary to make sure the plasma can besupported there, but as this is a local simulation this will ultimately have an influence the results."
+ For the top and bottom, For the top and bottom
+analvsis of full-disk magnetograms of the Michelson Doppler Luager aboard the Solar ancl lleliospherie Observatory (MDI/SOIIO) (Scherrer et al.,analysis of full-disk magnetograms of the Michelson Doppler Imager aboard the Solar and Heliospheric Observatory (MDI/SOHO) (Scherrer et al.
+ 1995)., 1995).
+ By adopting the clilferent detection. algorithm. and approaches. Meunier (2003) revealed. correlation of the network element number (or flix) will sunspots: in contrast. ILagenaar et al. (," By adopting the different detection algorithm and approaches, Meunier (2003) revealed correlation of the network element number (or flux) with sunspots; in contrast, Hagenaar et al. ("
+2003) declared some weak anti-correlated emergence rate of ERs and an independence of the total absolute flux for smaller network concentrations.,2003) declared some weak anti-correlated emergence rate of ERs and an independence of the total absolute flux for smaller network concentrations.
+ This discrepancy should be clarified wilh new analvsis., This discrepancy should be clarified with new analysis.
+ To clarify the problem and to close the debates are an essential task in understanding the solar evele phenomena., To clarify the problem and to close the debates are an essential task in understanding the solar cycle phenomena.
+ Fortunately. now MDI/SOLIO is providing a unique database - the full-disk magnetogranis over more (han 13 vears. covering the complete 23rd Solar Cycle.," Fortunately, now MDI/SOHO is providing a unique database - the full-disk magnetograms over more than 13 years, covering the complete 23rd Solar Cycle."
+ The 13.5 vear 5-min average full disk magnetograms are used in the current study., The 13.5 year 5-min average full disk magnetograms are used in the current study.
+ However. the poor temporal resolution makes (he identity of ERs questionable and the sensitivity of the full-disk magnetograms rules oul the possibility (to resolve the IN elements.," However, the poor temporal resolution makes the identity of ERs questionable and the sensitivity of the full-disk magnetograms rules out the possibility to resolve the IN elements."
+ Therefore. what we have identified in this study is basically the network magnetic elements.," Therefore, what we have identified in this study is basically the network magnetic elements."
+ In (his paper. we aim at learning the evelic variations of quiet Suns magnete flux aud small-scale magnetic elements.," In this paper, we aim at learning the cyclic variations of quiet Sun's magnetic flux and small-scale magnetic elements."
+ To use (he full-dise MDI magnetograms with the temporal coverage of entire Cvcle 23 comes from an awareness of (he intermittency of solar. cyclic behavior in both (he temporal aud spatial domains., To use the full-disc MDI magnetograms with the temporal coverage of entire Cycle 23 comes from an awareness of the intermittency of solar cyclic behavior in both the temporal and spatial domains.
+ By the intermittency to select. the maeneltogranms ol a short interval. e.g.. 10-30 hours. in a month lor each year. at the supposed. different. evcle phases.," By the intermittency to select the magnetograms of a short interval, e.g., 10-30 hours, in a month for each year, at the `supposed' different cycle phases."
+ would not guarantee a grape of the kev characteristics of a solar evele., would not guarantee a grape of the key characteristics of a solar cycle.
+ From our understandiug. to choose the database that cover the entire evele 23 is of overwhelming importance.," From our understanding, to choose the database that cover the entire cycle 23 is of overwhelming importance."
+ The database for the current study is unique in the sense that il is (he only space-borne magnetic measurements of the full Sun. for which the consistency in sensitivity and resolution persisted for a evele-long interval.," The database for the current study is unique in the sense that it is the only space-borne magnetic measurements of the full Sun, for which the consistency in sensitivity and resolution persisted for a cycle-long interval."
+ As we are interested in the elobal behavior of small-scale magnetic elements. saanpling network elements in a ¢eveleloue temporal domain and in all different. [lux ranges (or strengths) are more important than selecting a lew high cadence sequences interruptedly.," As we are interested in the global behavior of small-scale magnetic elements, sampling network elements in a cycle-long temporal domain and in all different flux ranges (or strengths) are more important than selecting a few high cadence sequences interruptedly."
+ Moreover. (he magnetic elements with different flux (or size) may have different origins ancl characteristics. therefore we group all the network magnetic elements into different. categories in accordance with their magnetic (his.," Moreover, the magnetic elements with different flux (or size) may have different origins and characteristics, therefore we group all the network magnetic elements into different categories in accordance with their magnetic flux."
+ In section 2. we describe (he observations. the technique of calibration. the evaluation of noise level of the maguelograms. the separation of active regions aud (he quiet Sun. ancl the selection of network elements.," In section 2, we describe the observations, the technique of calibration, the evaluation of noise level of the magnetograms, the separation of active regions and the quiet Sun, and the selection of network elements."
+ In section 3. we present the results of evelic behavior of quiet region magnetic flux and small-scale magnetic elements.," In section 3, we present the results of cyclic behavior of quiet region magnetic flux and small-scale magnetic elements."
+ Ii section 4. we make (he comparison wilh previous studies. and consider the possibilities on how (to unclerstand (the auti-correlated network magnetic elements with sunspolts.," In section 4, we make the comparison with previous studies, and consider the possibilities on how to understand the anti-correlated network magnetic elements with sunspots."
+ In section 5. we draw the conclusions.," In section 5, we draw the conclusions."
+Observations from the CAALA project (Driverefαἱ.2009:AJaldrvefal2009) are shown in Lig. 2..,"Observations from the GAMA project \citep{GAMA,Baldry09} are shown in Fig. \ref{GAMA obj}."
+ The source presented. has reported SDSS magnitudes of r(AB)=19.17 and. is found to be at redshifts of 220.276., The source presented has reported SDSS magnitudes of $r$ (AB)=19.17 and is found to be at redshifts of $z$ =0.276.
+ The spectrum is composed of ssec integrations. and. was processed with the use of 24 dedicated sky fibres to provide the master skv spectrum. for subtraction.," The spectrum is composed of $\times$ sec integrations, and was processed with the use of 24 dedicated sky fibres to provide the master sky spectrum for subtraction."
+ This master sky spectrum is shown scaled. to. for comparision., This master sky spectrum is shown scaled to for comparision.
+ While the sky subtraction does not hamper the identification of redshift for the source. there are systematic residuals in the data.," While the sky subtraction does not hamper the identification of redshift for the source, there are systematic residuals in the data."
+ The lower panel in Fig., The lower panel in Fig.
+ 2 shows the as defined in 2.4., \ref{GAMA obj} shows the as defined in \ref{skysubdefinition}.
+ The local error varies between close to the Poisson noise limit and over of the local skv spectrum wilh an rms 4.67+0.6 over the full spectrum., The local error varies between close to the Poisson noise limit and over of the local sky spectrum with an rms $\pm$ 0.6 over the full spectrum.
+ In the absence of a Poisson component from the target source continuum. the local error estimate is consistent with the variance estimation derived. from the pixel intensities during data reduction.," In the absence of a Poisson component from the target source continuum, the local error estimate is consistent with the variance estimation derived from the pixel intensities during data reduction."
+ The residual error histogram for the spectrum is shown in Fig. 3. ," The residual error histogram for the spectrum is shown in Fig. \ref{GAMA res1},"
+confirming that the local error estimate is consistent with the residual error in the source spectra which is in turn broadly consistent with the spectral error expected from the data reduction pixel variance., confirming that the local error estimate is consistent with the residual error in the source spectra which is in turn broadly consistent with the spectral error expected from the data reduction pixel variance.
+" Data from the WiggleZ survey (Drinkwaterοἱa£,2009). is in a cdillerent regime from that of the GAALA survey.", Data from the WiggleZ survey \citep{WiggleZ} is in a different regime from that of the GAMA survey.
+ Since WigeleZ sources are 20 « r(AD) < 22.5 high equivalent width texisci]] emission. line sources. there is. little source continuum to introduce increased noise to the error statistic., Since WiggleZ sources are 20 $<$ $r$ (AB) $<$ 22.5 high equivalent width ] emission line sources there is little source continuum to introduce increased noise to the error statistic.
+ Fig., Fig.
+ 4 shows such a WigeleZ survey source., \ref{WiggleZ obj} shows such a WiggleZ survey source.
+ Phe local error is close to the sky. limited case tthe limit expected from the pixel intensity. variance estiniates)., The local error is close to the sky limited case the limit expected from the pixel intensity variance estimates).
+ Since WigeleZ spectra are largely. [ree of continuum a stacking analvsis of the all 352 science fibre spectrain this frame should. produce an improvement in the local error estimate which scales as YN for the number of fibres stacked., Since WiggleZ spectra are largely free of continuum a stacking analysis of the all 352 science fibre spectrain this frame should produce an improvement in the local error estimate which scales as $\sqrt{\rm{N}}$ for the number of fibres stacked.
+ A clear departure from this scaling is seen in the upper panel of Vig. 5..," A clear departure from this scaling is seen in the upper panel of Fig. \ref{WiggleZ
+sky}."
+ his failure to follow the VN. suppression is the result of svstematic errors in the sky subtraction., This failure to follow the $\sqrt{\rm{N}}$ suppression is the result of systematic errors in the sky subtraction.
+ On applying the PCA correction to the sky subtracted spectra and repeating the stacking analwsis the local error estimate is seen to improve niarkedlv., On applying the PCA correction to the sky subtracted spectra and repeating the stacking analysis the local error estimate is seen to improve markedly.
+ This control of the build-up of svstematic error is easily seen in the residual error histogram of Fig. 6..," This control of the build-up of systematic error is easily seen in the residual error histogram of Fig. \ref{WiggleZ
+hist}."
+ Without the PCA correction the local error. estimator is clearly very dillerent from. the expected error. distribution (based. on pixel. variance estimates from data reduction)., Without the PCA correction the local error estimator is clearly very different from the expected error distribution (based on pixel variance estimates from data reduction).
+ With the application of the PCA correction prior to stacking. the local error estimator and variance data are largely in agreement.," With the application of the PCA correction prior to stacking, the local error estimator and variance data are largely in agreement."
+ To compare the previous results to observations taken with the nod-and-shullle observing technique we use deep observations taken. for quality control purposes. as part of the WigeleZ project.," To compare the previous results to observations taken with the nod-and-shuffle observing technique we use deep observations taken, for quality control purposes, as part of the WiggleZ project."
+ A 6hhour on source observation is available comprising three observations (taken over three nights) of 3.2400ssec., A hour on source observation is available comprising three observations (taken over three nights) of $\times$ sec.
+ The cata was observed in the Cross-Beam-Switching (CBS) mode with a pair of fibres allocated to cach target., The data was observed in the Cross-Beam-Switching (CBS) mode with a pair of fibres allocated to each target.
+ Phe observations where taken in the conventional A.XOmoega N|S mode with half of the available science fibres masked in order to provide the requisite CCD storage space (rather than the recently implemented. mini-shullling mode discussed later and in Sharp(2009))), The observations where taken in the conventional AAOmega N+S mode with half of the available science fibres masked in order to provide the requisite CCD storage space (rather than the recently implemented mini-shuffling mode discussed later and in \citet{minishuffle}) ).
+ Ao representative spectrum for a typical source (r(A3)222.26. 2=0.851) is presented. in Fig. 7..," A representative spectrum for a typical source $r$ (AB)=22.26, $z$ =0.851) is presented in Fig. \ref{wigglez ns3}."
+ The local error estimator indicates that the observations are indeed close to the Poisson noise limit expected. for the NS observing process., The local error estimator indicates that the observations are indeed close to the Poisson noise limit expected for the N+S observing process.
+ Observations of Taint sources require long duration exposures to build up the required signal-to-noise— ratio within cach spectrum., Observations of faint sources require long duration exposures to build up the required signal-to-noise ratio within each spectrum.
+ In the limit of pure Poisson-noise in the photon arrival rate. signal-to-noise is expected to build as the square-root of the observation time.," In the limit of pure Poisson-noise in the photon arrival rate, signal-to-noise is expected to build as the square-root of the observation time."
+ However. systematic defects in the spectrum. most commonly. from svstematic failure of the sky subtraction. will reduce the clliciency of this accrual of signal-to-noise.," However, systematic defects in the spectrum, most commonly from systematic failure of the sky subtraction, will reduce the efficiency of this accrual of signal-to-noise."
+ The common causes of these svsteniatie defects in fibre spectroscopy are considered. in appendix Al..., The common causes of these systematic defects in fibre spectroscopy are considered in appendix \ref{Appendix-problems}.
+ Determining the point at which each sky subtraction methodology becomes Limitect by systematic defects is Κον to selecting the most ellicient observational strategy to achieve ones goals., Determining the point at which each sky subtraction methodology becomes limited by systematic defects is key to selecting the most efficient observational strategy to achieve ones goals.
+ 1n the absence of an ideal data set we mimic long duration exposures using two methods., In the absence of an ideal data set we mimic long duration exposures using two methods.
+ First we consider the ellects of stacking Libre spectra from within a single MOS frame., First we consider the effects of stacking fibre spectra from within a single MOS frame.
+ Secondly we investigate stacking individual Libres [rom across multiple independent observations., Secondly we investigate stacking individual fibres from across multiple independent observations.
+ The former is not ideal as there may be strong correlated noise patterns between fibres in a single frame which would average out over many frames., The former is not ideal as there may be strong correlated noise patterns between fibres in a single frame which would average out over many frames.
+ Ehe later is dillicult in practise as à suitable data set. which places a single fibre on blank sky across many independent. observing frames. is not usually available.," The later is difficult in practise as a suitable data set, which places a single fibre on blank sky across many independent observing frames, is not usually available."
+ The residual error as a function. of time for a single [rame stacking process are shown in Fig. &.., The residual error as a function of time for a single frame stacking process are shown in Fig. \ref{noPCA stack}.
+ Single frames from the CAMA. WigeleZ and WigeleZ N18 cata sets presented previously are used.," Single frames from the GAMA, WiggleZ and WiggleZ N+S data sets presented previously are used."
+ Additionally. in order to construct a single fibre multi-frame stack we use the properties of the WigeleZ survey cata to construct a stack of | 200 individual fibre spectra from data taken throughout 2009 (Lig. 99).," Additionally, in order to construct a single fibre multi-frame stack we use the properties of the WiggleZ survey data to construct a stack of $+$ 200 individual fibre spectra from data taken throughout 2009 (Fig. \ref{PCA stack}) )."
+ In cach case stacked spectra are created with a range of values of N. the number of fibres used.," In each case stacked spectra are created with a range of values of N, the number of fibres used."
+ For cach value of N. 100 random realisations of the stacked spectrum. are created and the local error and its rims scatter are plotted against the number of spectra used in the stack.," For each value of N, 100 random realisations of the stacked spectrum are created and the local error and its rms scatter are plotted against the number of spectra used in the stack."
+ Each stack of N fibres simulates an exposure of hhours on blank sky., Each stack of N fibres simulates an exposure of hours on blank sky.
+ The rate of reduction in residual error with time is then compared. to the YN rate of decline expected. Crom pure Poisson statistics., The rate of reduction in residual error with time is then compared to the $\sqrt{\rm{N}}$ rate of decline expected from pure Poisson statistics.
+ Lt is clear from Fig., It is clear from Fig.
+ S. that for both the data sets which use purely the subtraction. a svstematic noise Hoor is reach after 10hhours.," \ref{noPCA stack} that for both the data sets which use purely the subtraction, a systematic noise floor is reach after $\sim$ hours."
+ With the application of the PCA procedure to the data prior to stacking there is, With the application of the PCA procedure to the data prior to stacking there is
+outflows sugeestOO that the eas behind the shock mueht contain a range of temperatures as well as regions of hieh density (O'Brienetal.1991: see also Llovdctal.1996)).,outflows suggest that the gas behind the shock might contain a range of temperatures as well as regions of high density \citealt{obrien94}; see also \citealt{lloyd96}) ).
+ The required gas densities. teniperatures. ane total masses are not excessive.," The required gas densities, temperatures, and total masses are not excessive."
+ Atcon 5 kpe. the observed SGCz eunission around dav 10 |ο produced by a shell with thickness ~107ean (outer radius ~5ν10 en). deusity ngo100n3.Lo teperature 105K. and total mass of a few times ONES.," At 5 kpc, the observed GHz emission around day 40 could be produced by a shell with thickness $\sim10^{13}\rm\,cm$ (outer radius $\sim5\times10^{14}\rm\,cm$ ), density $n_e\sim10^7\rm\,cm^{-3}$, temperature $\sim10^6\rm\,K$ , and total mass of a few times $10^{-6}\rm\,M_\sun$."
+ Whether such models can reproduce the observed Πο curves and spectra remains to be seen., Whether such models can reproduce the observed light curves and spectra remains to be seen.
+" Any imiodel iust also explain why thebump secu inκ... His not seen in the radio Πο curves of other classical uovae. although may classical novae show evidence of shocks,"," Any successful model must also explain why the bump seen in is not seen in the radio light curves of other classical novae, although many classical novae show evidence of shocks."
+" These models. and the relation between aad “other uovae. depend crucially on continued racio monitoring. which will show whether uultimatelv follows the canonical TubLeft""v models at late times."," These models, and the relation between and other novae, depend crucially on continued radio monitoring, which will show whether ultimately follows the canonical Hubble-flow models at late times."
+ Upcoming radio imaging with the EVLA in its most extended. highest resolution “A” configuration will also provide valuable information on the plysical sizeand eeometry of the ejecta.," Upcoming radio imaging with the EVLA in its most extended, highest resolution “A” configuration will also provide valuable information on the physical sizeand geometry of the ejecta."
+erows. as vost new lines occur at longer aud longer waveleueths (bulk modes).,"grows, as most new lines occur at longer and longer wavelengths (bulk modes)."
+ Also. this is hardly applicable to the trio structures.," Also, this is hardly applicable to the trio structures."
+ A more cjificient wav consists in slightly modifving a eiven strucure. by acing a small eroup. e.g. methyl or bydroxvl. as was done for trios in 1.," A more efficient way consists in slightly modifying a given structure, by displacing a small group, e.g. methyl or hydroxyl, as was done for trios in Fig."
+ , 4.
+Bu the nuniber of alternatives is rather lanited for simall strctures., But the number of alternatives is rather limited for small structures.
+ he otjer hand. evidence fro1i a broad range of meteorites and Organe er in Comet SIP/Wild 2 samples does not point to abunudau. large. PAIIs iusead. if points to a svuthlesis of small reactive molecules. whose raloni condensation and subsequent reamransement chemistry leads to a luehIv cross-luked macromolecule (see Cody et al. (2008))).," On the other hand, evidence from a broad range of meteorites and organic matter in Comet 81P/Wild 2 samples does not point to abundant, large, PAHs; instead, it points to a synthesis of small reactive molecules, whose random condensation and subsequent rearrangement chemistry leads to a highly cross-linked macromolecule (see Cody et al. \cite{cod}) )."
+ This is precisely the nodel upon which new structures were created below., This is precisely the model upon which new structures were created below.
+ Iu order to produce composie structures. the eClicary structures of Sec.," In order to produce composite structures, the elementary structures of Sec."
+ 2 can be associated either by ightly counectiug them. with short aud multiple ids. or by concatenatiue f1011 end to cud wit1 siuee bonds or small clains.," 2 can be associated either by tightly connecting them with short and multiple bonds, or by concatenating them end to end with single bonds or small chains."
+ arge nuniber of structures were built along these lines., A large number of structures were built along these lines.
+ Lack of sp:ice does allow the inclusion of a] illustrative sketches: oily the leading member each retaimed family is 1lustrated iu the followiic figures., Lack of space does not allow the inclusion of all illustrative sketches; only the leading member of each retained family is illustrated in the following figures.
+ Towever. the owlng accolmpanving text is intended to stress tha an iucreasime nunber of inor structural modificatiois of aptly selected compsite structures. made up of he few molecules described in Sec.," However, the following accompanying text is intended to stress that an increasing number of minor structural modifications of aptly selected composite structures, made up of the few molecules described in Sec."
+ 2. and combiner in the right proportions. will produce increasingly dense concentratious of liies in the right spectral bans.," 2, and combined in the right proportions, will produce increasingly dense concentrations of lines in the right spectral bands."
+ Figure 7 shows a sample of this type: compacted trios (a). made of 2 nitrogen trios (sec fig.," Figure 7 shows a sample of this type: compacted trios (a), made of 2 nitrogen trios (see fig."
+ 1) tightly connected through their utiech atoms (19 at.)., 4) tightly connected through their nitrogen atoms (49 at.).
+ The other samples to be used (but not shown here) are: bj sue as (a). except for the utrogen atous beiie substituted with C atoms and the attendant creation of wo C=C bonds {7 at.):," The other samples to be used (but not shown here) are: b) same as (a), except for the nitrogen atoms being substituted with C atoms and the attendant creation of two C=C bonds (47 at.);"
+ ο) structures (a) and (b) tightv connected (LOL at.), c) structures (a) and (b) tightly connected (104 at.);
+ d) structure (0) plus 1 authrac‘one and 1 phenawrene (169 at.):, d) structure (c) plus 1 anthracene and 1 phenantrene (169 at.);
+ ο) same as (d). except for auoher two short CIT» links (190 at.).," e) same as (d), except for another two short $_{2}$ links (190 at.)."
+ The spectra of these 5 structiwes are concatenated mto the lower spectra of Fie., The spectra of these 5 structures are concatenated into the lower spectrum of Fig.
+ 11., 11.
+ As expected. the imiuber of active ir lines notably increased as conrpared with isolated trios (upper spectrum. Fie.," As expected, the number of active ir lines notably increased as compared with isolated trios (upper spectrum, Fig."
+ 5)., 5).
+ The spectral density, The spectral density
+"of Figure 8 are the PMS isochrones of for ages of 1, 2, 4, 8, 16, 32 and 64 Myr.","of Figure \ref{fig_hist} are the PMS isochrones of \cite{sie00} for ages of 1, 2, 4, 8, 16, 32 and 64 Myr."
+ Note that the photometric uncertainty of our data (see large crosses on the figure) confirms that it is possible to assign relative ages to these stars with an accuracy of better than factor of 2., Note that the photometric uncertainty of our data (see large crosses on the figure) confirms that it is possible to assign relative ages to these stars with an accuracy of better than a factor of 2.
+ Solid thick points and asterisks in the figure representa all bona-fide PMS objects (i.e. those with Ho excess emission) having ages in the range 1—64 MMyr and masses between 0.9 and 4MMo., Solid thick points and asterisks in the figure represent all bona-fide PMS objects (i.e. those with $\alpha$ excess emission) having ages in the range $1 - 64$ Myr and masses between $0.9$ and $_\odot$.
+ The corresponding age distribution is shown by the solid histogram in the right panel of the figure., The corresponding age distribution is shown by the solid histogram in the right panel of the figure.
+" In the same we also show, as a dashed line, the of the age paneldistribution of all stars with no Ha excess histogramyounger than MMyr and with masses in the range from 0.9 and 4MMo, marked as small dots in the left-hand panel (note that the few bona-fide PMS stars outside of these and mass ranges, remainingshown as thick grey points, and are not ageconsidered here)."," In the same panel we also show, as a dashed line, the histogram of the age distribution of all stars with no $\alpha$ excess younger than Myr and with masses in the range from $0.9$ and $_\odot$, marked as small dots in the left-hand panel (note that the few remaining bona-fide PMS stars outside of these age and mass ranges, shown as thick grey points, and are not considered here)."
+" For the objects with no Ho excess emission we have only sampled the age distribution younger than MMyr, in order to avoid the significant field star contamination in the bluest part of the CMD."," For the objects with no $\alpha$ excess emission we have only sampled the age distribution younger than Myr, in order to avoid the significant field star contamination in the bluest part of the CMD."
+" For comparison purposes, their histogram has been normalized vertically (i.e. shifted by 1.1 dex) so as to match the distribution of the bona-fide PMS stars at the youngest age."," For comparison purposes, their histogram has been normalized vertically (i.e. shifted by $1.1$ dex) so as to match the distribution of the bona-fide PMS stars at the youngest age."
+" The histograms of reffig,istshowthatstar formationinNGC 33603hasbeenongoing f —20Myrandnogapsareevident, atleastatthelevelo f bOnl)eage(a f actoro ftwo,asshownbythesizeo f thebins)."," The histograms of \\ref{fig_hist} show that star formation in 3603 has been ongoing for at least 10--20 Myr and no gaps are evident, at least at the level of resolution that we have adopted for the age (a factor of two, as shown by the size of the bins)."
+"I f excess emission, for which no selection effects are present other than photometric completeness that is nonetheless always >85%, we would have to conclude that over the ~16 the star formation rate in this has pastbeen progressivelyMMyr increasing."," If we consider stars with no $H\alpha$ excess emission, for which no selection effects are present other than photometric completeness that is nonetheless always $> 85\,\%$, we would have to conclude that over the past $\sim 16$ Myr the star formation rate in this region has been progressively increasing."
+" However, it is importantregion to consider that many of the older stars might have migrated out of our FoV. According to Rochauetal.(2010),, the velocity dispersion of stars in the central regions of 33603 is ~4.5 km s! and appears to be pretty constant in the mass range that they sample (1.7—9M5)."," However, it is important to consider that many of the older stars might have migrated out of our FoV. According to \cite{roc10}, the velocity dispersion of stars in the central regions of 3603 is $\sim 4.5$ km $^{-1}$ and appears to be pretty constant in the mass range that they sample $\sim1.7-9M_{\odot}$ )."
+" This implies that a 10MMyr old star would have had time to travel as far as ppc away from its birthplace, i.e. well beyond the ~5x ppc? area covered by our observations."," This implies that a Myr old star would have had time to travel as far as pc away from its birthplace, i.e. well beyond the $\sim 5 \times 5$ $^2$ area covered by our observations."
+" This might be one of the causes of the observed drop in the number of stars with increasing age shown in the histogram of reffig,ist, andwouldalsoexplainthesomewhatdif ferent radialdist ribu"," This might be one of the causes of the observed drop in the number of stars with increasing age shown in the histogram of \\ref{fig_hist}, and would also explain the somewhat different radial distributions of old and young PMS stars seen in Figure \ref{fig_ks}."
+tioi , A survey of a wider area around 3603 is needed to properly address the evolution of the star formation rate in this cluster.
+"Interestingly, the age distribution of bona-fide PMS stars does not seem to differ in any systematic way within the error bars from that of objects with no Ha excess, seemingly suggesting that the ratio of PMS stars with and without Ha excess emission is not a function of age."," Interestingly, the age distribution of bona-fide PMS stars does not seem to differ in any systematic way within the error bars from that of objects with no $\alpha$ excess, seemingly suggesting that the ratio of PMS stars with and without $\alpha$ excess emission is not a function of age."
+" This might appear at odds with common wisdom suggesting that the efficiency of the accretion process at the origin of the Ha emission should decrease with time as a PMS objects approaches the MS (seeSicilia-Aguilaratal.2010;DeMarchiet2010a,andreferences therein).."," This might appear at odds with common wisdom suggesting that the efficiency of the accretion process at the origin of the $\alpha$ emission should decrease with time as a PMS objects approaches the MS \citep[see][and references therein]{sic10,DeM10}."
+" However, selection effects here can be important."," However, selection effects here can be important."
+" In particular, as PMS objects approach the MS, both their Ha and bolometric luminosities decrease, but not necessarily in the same way."," In particular, as PMS objects approach the MS, both their $\alpha$ and bolometric luminosities decrease, but not necessarily in the same way."
+" If the bolometric luminosity drops more rapidly than the Ha luminosity, PMS stars of older ages will become easier to identify for our method since it orattpusesQa excess emission at the 5c level or higher (seeDeresolutionthatwkhaneudspl"," If the bolometric luminosity drops more rapidly than the $\alpha$ luminosity, PMS stars of older ages will become easier to identify for our method since it requires an excess emission at the $5\,\sigma$ level or higher \citep[see][]{DeM10}."
+"ed This effect can have important implications on our understanding of the accretion process and of the star formation rate, which we will address in a forthcoming the star formation process in a number of paperyoung comparingclusters."," This effect can have important implications on our understanding of the accretion process and of the star formation rate, which we will address in a forthcoming paper comparing the star formation process in a number of young clusters."
+" We have already started to study the stellar population in the DDoradus region in the LMC, applying the same observational strategy described in this"," We have already started to study the stellar population in the Doradus region in the LMC, applying the same observational strategy described in this"
+where e(r) is the wind velocity at cach point in the heating shell from the time-dependent solution.,where $v(r)$ is the wind velocity at each point in the heating shell from the time-dependent solution.
+ “Phe problem with this is that at the inner edge of the heating shell the heating rate is finite while the velocity is very close to zero. resulting in a slight overestimate of the total energy input near the inner edge of the shell in the steady wind solution.," The problem with this is that at the inner edge of the heating shell the heating rate is finite while the velocity is very close to zero, resulting in a slight overestimate of the total energy input near the inner edge of the shell in the steady wind solution."
+ Care must also be taken in integrating through the singular point in equation (11)) at A4?=1., Care must also be taken in integrating through the singular point in equation \ref{eq:machno}) ) at $M^2=1$.
+ Most authors (e.g. Lamers&Cassinelli 1999)) solve the steacky wind equations starting from this point but for our purposes it is better to start the integration outside of, Most authors (e.g. \citealt{lc99}) ) solve the steady wind equations starting from this point but for our purposes it is better to start the integration outside of
+where the true value of Lf.) 15 that which makes the right hand side of (10) equal to unity.,where the true value of $H_{exp}$ is that which makes the right hand side of (10) equal to unity.
+ For convenience we shall refer to the statistic represented by (10) as Y(LLeep)., For convenience we shall refer to the statistic represented by (10) as $Y(H_{exp})$.
+" Lastly from Fig 1 and equation (7) where //,,,, is the value obtained above.", Lastly from Fig 1 and equation (7) where $H_{exp}$ is the value obtained above.
+" The above analvsis applies to anv one region undergoing Hubble-like expansion at a rate IH,Jp:", The above analysis applies to any one region undergoing Hubble-like expansion at a rate $H_{exp}$.
+ The region with the sub-IIubble expansion rate (//.) found by TOS lies between A2. aud lHipsidnFie.l. (ο.," The region with the sub-Hubble expansion rate $H_{v}$ ) found by T08 lies between $R_{v}$ and $R_{ES}$ in Fig.1. $R_{v}$,"
+" Res and I1, ave defined in the next section).", $R_{ES}$ and $H_{v}$ are defined in the next section).
+" Devond py the expansion is assumed to return to the elobal Lhibble rate Z£, (2(2)). (", Beyond $R_{ES}$ the expansion is assumed to return to the global Hubble rate $H_{o}$ $H(z)$ ). (
+At higher z the global rate becomes Πιτ}=Πω]—Ov)4z)*+ QíJE7).,"At higher $z$ the global rate becomes $H(z)=H_{o}
+((1-\Omega_{\Lambda})(1+z)^3+\Omega_{\Lambda})^{1/2}$ )."
+" Similarly. by (11) in each of these regions we then have FI.) aad RT(,)."," Similarly, by (11) in each of these regions we then have $R(H_{v})$ and $R(H(z_{g}))$."
+" With (wo expanding regions. (he four kev quantities in the analvsis are (he mean of the ralio y'/i' (equ'n 10) evaluated with 77, and H(z74) replacing 77,,,, designated here Y(/7,.) andY(1I(z,)) and the normalized distances ΠΕΠ)B. and B(IT(z))/H,."," With two expanding regions, the four key quantities in the analysis are the mean of the ratio $y^{\prime}/x^{\prime}$ (equ'n 10) evaluated with $H_{v}$ and $H(z_{g})$ replacing $H_{exp}$ designated here $Y(H_{v})$ and$Y(H(z_{g}))$ and the normalized distances $R(H_{v})/R_{v}$ and $R(H(z_{g}))/R_{v}$."
+ These four quantities are calculated for every field galaxy., These four quantities are calculated for every field galaxy.
+" In order to confirm the TOS result. Y(77,) should be unity when |<A(GT)/B.<Res/B lor the particular value of H,,,=IT, and Y(11(2,))πα) should ΠΕ Πο "," In order to confirm the T08 result, $Y(H_{v})$ should be unity when $1 \leq R(H_{v})/R_{v} \leq R_{ES}/R_{v}$ for the particular value of $H_{exp}=H_{v}$ and $Y(H(z_{g}))$ should be unity when $R(H(z_{g}))/R_{v}$ becomes larger than $R_{ES}/R_{v}$."
+Ου rely on the assumption that the galaxies are randomly distributed about the group centers for the sample as a whole., Both of the above conditions rely on the assumption that the galaxies are randomly distributed about the group centers for the sample as a whole.
+ Note also that because we are dealing with ratios. the result is independent of the actual value assumed [ον 77.," Note also that because we are dealing with ratios, the result is independent of the actual value assumed for $H_{o}$."
+ In order to validate the above procedure a simple numerical simulation was performed., In order to validate the above procedure a simple numerical simulation was performed.
+ In (his simulation the same group properties i.e. (he positions. redshifts and masses (Irom which τν could be calculated) of groups used in (he analvsis of the observations were used as centers around which [field galaxies’ were randomly distributed (both in angle 9 aud distance AR).," In this simulation the same group properties i.e. the positions, redshifts and masses (from which $R_{v}$ could be calculated) of groups used in the analysis of the observations were used as centers around which `field galaxies' were randomly distributed (both in angle $\theta$ and group-centric distance $R$ )."
+" Around each of the 12620 group centers are 300 randomly chosen field galaxies,", Around each of the 12620 group centers are 300 randomly chosen field galaxies.
+ For a given cosmology. line-o[-sight. redshiifts were assigned (either as a result. of elobal IIubble expansion or sub- Hubble expansion about the group center) depending on whether τν was greater (han or less than 2.," For a given cosmology, line-of-sight redshifts were assigned (either as a result of global Hubble expansion or sub- Hubble expansion about the group center) depending on whether $R/R_{v}$ was greater than or less than 2."
+ The simulation is clearly idealized to the extent (hat each eroup and its retinue of randomly distributed galaxies was (reated as an isolated svstem., The simulation is clearly idealized to the extent that each group and its retinue of randomly distributed galaxies was treated as an isolated system.
+ The data were (hen analvsed (as described above) using the sameroutines that were used for the observations., The data were then analysed (as described above) using the sameroutines that were used for the observations.
+ Here (he values of 2 plotted. on the abscissa come from the simulation., Here the values of $R$ plotted on the abscissa come from the simulation.
+ Shown in Fig., Shown in Fig.
+" 2 are the results for an assumed. (0,,.04. 05) (0.3.0.7 0.0) cosmology."," 2 are the results for an assumed $\Omega_{m},\Omega_{\Lambda},\Omega_{k}$ ) (0.3,0.7 ,0.0) cosmology."
+" The red points represent Y(/7,) and R/ BI.", The red points represent $Y(H_{v})$ and $R/R_{v}$ .
+" The black points were calculated assuming LT,=IT(z,) and plotted as Y'(/7(2,)) versus H/ Ii.", The black points were calculated assuming $H_{exp}=H(z_{g})$ and plotted as $Y(H(z_{g}))$ versus $R/R_{v}$ .
+ Note the characteristic step in, Note the characteristic step in
+"Stot=2.96+0.51 Jy km s! (?) and Si,=3.07+0.36 Jy km s! (?) (this latter value is without the correction for self absorption).",$S_{\rm{tot}}=2.96 \pm 0.51$ Jy km $^{-1}$ \citep{Matthews2001b} and $S_{\rm{tot}}=3.07 \pm 0.36$ Jy km $^{-1}$ \citep{Springob2005} (this latter value is without the correction for self absorption).
+" masses were estimated using the luminosity distance Dry, given by: We assumed Hy=70 km sec! Mpc!, Qm=0.27 and Qa=0.73 (qo= -0.50)."," masses were estimated using the luminosity distance $D_{\rm{L}}$ , given by: We assumed $H_{0}=70$ km $^{-1}$ $^{-1}$, $\Omega_{\rm{m}}=0.27$ and $\Omega_{\rm{\Lambda}}=0.73 \,$ $q_{0}=-0.59$ )."
+" Since we are dealing with low redshift galaxies, the exact value of go affects our results only slightly."," Since we are dealing with low redshift galaxies, the exact value of $q_{0}$ affects our results only slightly."
+" The redshift was calculated as z=V;y,/c, where Vyys is the galaxy systemic velocity derived from the tilted ring analysis (see section 2.2))."," The redshift was calculated as $z=V_{\rm{sys}}/c$ , where $V_{\rm{sys}}$ is the galaxy systemic velocity derived from the tilted ring analysis (see section \ref{RC}) )."
+ Our results are listed in table 1.., Our results are listed in table \ref{tab:HIres}.
+" The derivation of the velocity fields is not straightforward because of beam-smearing effects, which are particularly severe here because of the low spatial resolution."," The derivation of the velocity fields is not straightforward because of beam-smearing effects, which are particularly severe here because of the low spatial resolution."
+" Indeed, especially in the inner regions, the 21-cm line profiles are unusually broad and asymmetric."," Indeed, especially in the inner regions, the 21-cm line profiles are unusually broad and asymmetric."
+" A Gaussian fit or an intensity-weighted mean (IWM), as used by ?,, give velocities far from the profile peaks and biased towards the systemic velocity."," A Gaussian fit or an intensity-weighted mean (IWM), as used by \citet{Pickering1997}, give velocities far from the profile peaks and biased towards the systemic velocity."
+ This leads to velocity fields with systematically underestimated inner velocity gradients and consequently to slowly rising rotation curves., This leads to velocity fields with systematically underestimated inner velocity gradients and consequently to slowly rising rotation curves.
+" This is illustrated by figure 7 of ?,, which shows the rotation curves of Malin 1 and of NGC 7589 overlaid on slices of the datacubes along the major axes of the galaxies."," This is illustrated by figure 7 of \citet{Pickering1997}, which shows the rotation curves of Malin 1 and of NGC 7589 overlaid on slices of the datacubes along the major axes of the galaxies."
+" Clearly, the inner points of the rotation curves are at lower velocities than the emission peaks, suggesting that they are underestimated."," Clearly, the inner points of the rotation curves are at lower velocities than the emission peaks, suggesting that they are underestimated."
+ The importance of beam-smearing effects is demonstrated by the galaxy models presented in section 2.4.., The importance of beam-smearing effects is demonstrated by the galaxy models presented in section \ref{sub:models}.
+ We have derived new velocity fields with a different technique to minimize beam-smearing effects., We have derived new velocity fields with a different technique to minimize beam-smearing effects.
+" We selected velocities near the peaks of the emission, estimating an IWM velocity from the upper part of the line profile (above ~75% of the peak intensity) and neglecting the broad wings."," We selected velocities near the peaks of the emission, estimating an IWM velocity from the upper part of the line profile (above $\sim 75\%$ of the peak intensity) and neglecting the broad wings."
+" This is like taking the profile peak velocities, but with the advantage of less noise in the resulting velocity field."," This is like taking the profile peak velocities, but with the advantage of less noise in the resulting velocity field."
+" In the outer regions of the galaxies, where the line profiles are more regular and symmetric, we used a Gaussian fit."," In the outer regions of the galaxies, where the line profiles are more regular and symmetric, we used a Gaussian fit."
+ The resulting velocity fields are shown in figure 1.., The resulting velocity fields are shown in figure \ref{fig:Maps}.
+" The left panels of figure 1 show the R-band image of Malin 1 from ? (top), the total map (middle) and the velocity field (bottom)."," The left panels of figure \ref{fig:Maps} show the R-band image of Malin 1 from \citet{MooreParker2006} (top), the total map (middle) and the velocity field (bottom)."
+ The distribution is approximately as extended as the optical LSB disk., The distribution is approximately as extended as the optical LSB disk.
+" The outer radius measures ~64"", which corresponds to ~100 kpc."," The outer radius measures $\sim 64''$, which corresponds to $\sim 100$ kpc."
+" On the southern side, a spiral arm is visible."," On the southern side, a spiral arm is visible."
+ It extends for more than 80 kpc and apparently it is the gaseous counterpart of the spiral arm detected by ?.., It extends for more than 80 kpc and apparently it is the gaseous counterpart of the spiral arm detected by \citet{MooreParker2006}.
+" The velocity field shows a strong variation of the position angle of the major axis with radius, both in the approaching and in the receding sides of the galaxy."," The velocity field shows a strong variation of the position angle of the major axis with radius, both in the approaching and in the receding sides of the galaxy."
+ Such behaviour is usually attributed to the presence of a warp., Such behaviour is usually attributed to the presence of a warp.
+" Overall, the velocity field is regular and symmetric."," Overall, the velocity field is regular and symmetric."
+" The right panels of figure | show the R-band image of NGC 7589 from ? (top), the total map (middle) and the velocity field (bottom)."," The right panels of figure \ref{fig:Maps} show the R-band image of NGC 7589 from \citet{Galaz2006} (top), the total map (middle) and the velocity field (bottom)."
+ The distribution is as extended as the optical LSB disk., The distribution is as extended as the optical LSB disk.
+" The outer radius measures ~80"", which corresponds to ~48 kpc."," The outer radius measures $\sim 80''$, which corresponds to $\sim 48$ kpc."
+" The distribution follows that of the light, with two main concentrations (North-East and South-West) and two faint outer spiral arms."," The distribution follows that of the light, with two main concentrations (North-East and South-West) and two faint outer spiral arms."
+ The velocity field is regular and symmetric., The velocity field is regular and symmetric.
+ Kinematical parameters and rotation curves were derived by fitting a tilted-ring model to the observed velocity fields (?).., Kinematical parameters and rotation curves were derived by fitting a tilted-ring model to the observed velocity fields \citep{Begeman1987}.
+ We assumed purely circular motions and described the galaxy by a set of concentric rings., We assumed purely circular motions and described the galaxy by a set of concentric rings.
+ We used all the points of the velocity fields to maximize the statistics of the least-square fit., We used all the points of the velocity fields to maximize the statistics of the least-square fit.
+" The points were weighted by cos?(6), where 0 is the azimuthal angle in the plane of the galaxy."," The points were weighted by $cos^{2}(\theta)$, where $\theta$ is the azimuthal angle in the plane of the galaxy."
+" As a first step, we determined the kinematical centre (Qin.Okin) and the mean systemic velocity V;,,."," As a first step, we determined the kinematical centre $(\alpha_{\rm{kin}}, \delta_{\rm{kin}})$ and the mean systemic velocity $V_{\rm{sys}}$."
+" Then, we Κερί (Qin,Okin) and Vey, fixed and determined the position angle P.A. as a function of radius."," Then, we kept $(\alpha_{\rm{kin}}, \delta_{\rm{kin}})$ and $V_{\rm{sys}}$ fixed and determined the position angle $P.A.$ as a function of radius."
+" Subsequently, we tried to determine the inclination angle i."," Subsequently, we tried to determine the inclination angle $i$."
+" Unfortunately, because of the low spatial resolutions, inclination angles and rotational velocities are strongly correlated and no independent determination of i was possible."," Unfortunately, because of the low spatial resolutions, inclination angles and rotational velocities are strongly correlated and no independent determination of $i$ was possible."
+" Therefore, we kept the inclination angles constant, adopting values derived from optical observations."," Therefore, we kept the inclination angles constant, adopting values derived from optical observations."
+ We used i=38? for Malin 1 (?) and i=58? for NGC 7589 (?).., We used $i = 38^{\circ}$ for Malin 1 \citep{MooreParker2006} and $i = 58^{\circ}$ for NGC 7589 \citep{Pickering1997}.
+ Inclinations derived from the total maps are consistent with those values., Inclinations derived from the total maps are consistent with those values.
+" Finally, we determined the circular velocity V. in each ring,keeping all the other parameters fixed."," Finally, we determined the circular velocity $V_{\rm{c}}$ in each ring,keeping all the other parameters fixed."
+ Our results are listed in table 2 and 3.., Our results are listed in table \ref{tab:kinparams} and \ref{table:RC}.
+" In the tilted-ring analysis of NGC 7589, we used a ring width of 20""—1 beam."," In the tilted-ring analysis of NGC 7589, we used a ring width of $20'' \sim 1$ beam."
+" For the kinematical centre, we obtained values which are consistent within the errors with the optical ones."," For the kinematical centre, we obtained values which are consistent within the errors with the optical ones."
+" The values obtained for V,,, and P.A. approximately agree with those of ?..", The values obtained for $V_{\rm{sys}}$ and $P.A.$ approximately agree with those of \citet{Pickering1997}. .
+" For Malin 1, we used a ring width of 10""—1/2 beam."," For Malin 1, we used a ring width of $10'' \sim 1/2$ beam."
+" To determine the kinematical centre, we used the points inside 40""—62 kpc."," To determine the kinematical centre, we used the points inside $40'' \sim 62$ kpc."
+ Thederived values agree within the errors with the optical ones., Thederived values agree within the errors with the optical ones.
+" The value of V,,, agrees within the errors with that derived by ?.. The P.A. for Malin 1 isnot constant.", The value of $V_{\rm{sys}}$ agrees within the errors with that derived by \citet{Pickering1997}.. The $P.A.$ for Malin 1 isnot constant.
+" For theinner 15""~24 kpc we found —1?.0+1°.6 and for the outer radii a linear increase up to 65°.1+4°.9 at R=65” (see the", For theinner $15'' \sim 24$ kpc we found $-1^{\circ}.0 \pm 1^{\circ}.6$ and for the outer radii a linear increase up to $65^{\circ}.1 \pm 4^{\circ}.9$ at $R = 65''$ (see the
+star formation listorics.,star formation histories.
+ No attempt was made to include galaxy evolution effects. which may result in elliptical galaxy colors that are sliebitlv too red at high redshift (sec. e.g. Druzual |1996]).," No attempt was made to include galaxy evolution effects, which may result in elliptical galaxy colors that are slightly too red at high redshift (see, e.g., Bruzual [1996])."
+ The resulting galaxy colors aeree with those computed bv Fukugita. Shiuasaku. Ichikawa (1995) to within O.liuas for most filters. galaxy spectral types. aud redshifts: this is acceptable eiveu the differences in the input spectra in the two works.," The resulting galaxy colors agree with those computed by Fukugita, Shimasaku, Ichikawa (1995) to within $0.1 \mag$ for most filters, galaxy spectral types, and redshifts; this is acceptable given the differences in the input spectra in the two works."
+ Galaxy spectra at optical waveleueths are a superposition of diverse stellar spectra. modified slightly bv nebulu ciission aud dust effects," Galaxy spectra at optical wavelengths are a superposition of diverse stellar spectra, modified slightly by nebular emission and dust effects."
+ This superposition implies that ogalaxies have ess sharply peaked spectra than individual stars and so lie on the same side of the color-color space stellar locus as power law spectra., This superposition implies that galaxies have less sharply peaked spectra than individual stars and so lie on the same side of the color-color space stellar locus as power law spectra.
+ Tadeccl. he spectruni of a star forming ealaxy can be a good power law over ost of the optical window.," Indeed, the spectrum of a star forming galaxy can be a good power law over most of the optical window."
+" Fortunately, the 10004. break due to stellar atinosplierie features interrupts lis approxiuate power lav and can be used o distinguish. almost all galaxies at :=.1 roni afterglow candidates eiven sufficieut color information."," Fortunately, the $4000$ break due to stellar atmospheric features interrupts this approximate power law and can be used to distinguish almost all galaxies at $z \la 1$ from afterglow candidates given sufficient color information."
+ With observatious iu three filters. here is some redshift where the 10004. ybreak lies in the ολο filter. aud galaxy colors can match those of a pure power law.," With observations in three filters, there is some redshift where the $4000$ break lies in the middle filter, and galaxy colors can match those of a pure power law."
+ The addition of a fourth filter breaks this degeueracv., The addition of a fourth filter breaks this degeneracy.
+ For example. im he B.RI plane. an nreeular ealaxy at zmθ lies on the power law sequence: however. its BR color is imuuchn redder than the correspouding power law.," For example, in the $U-B$, $R-I$ plane, an irregular galaxy at $z \approx 0.4$ lies on the power law sequence; however, its $B-R$ color is much redder than the corresponding power law."
+ In addition to color criteria. it is of course possible to distinguis[um ealaxies from CRB afterelows morphologically.," In addition to color criteria, it is of course possible to distinguish galaxies from GRB afterglows morphologically."
+ The fraction of galaxies that appear extended depends on the available spatial resolution. with the a few ealaxics appearing poiutlike even 1- IIubble Space Telescope images. bu subarceseconcd resolution will iu practice greatly reduce confusio- due to galaxies.," The fraction of galaxies that appear extended depends on the available spatial resolution, with the a few galaxies appearing pointlike even in Hubble Space Telescope images, but subarcsecond resolution will in practice greatly reduce confusion due to galaxies."
+ Quasar colors were calculated using the composite quasar specruni by Francis et al (1991). extrapolated to AD0.640008 with fXxA 248) ," Quasar colors were calculated using the composite quasar spectrum by Francis et al (1991), extrapolated to $\lambda > 0.6\micron$ with $f_\lambda \propto \lambda^{-2.18}$ ."
+This extrapolation will result in nodest color errors for low redshits where the Πα line falls in a iucasured filter. hit should not affect our conclusions greatly.," This extrapolation will result in modest color errors for low redshifts where the $\alpha$ line falls in a measured filter, but should not affect our conclusions greatly."
+ Quasar spectra are reasonably well approximated by power laws (plus broad cussion lines). across a wide waveleneth range., Quasar spectra are reasonably well approximated by power laws (plus broad emission lines) across a wide wavelength range.
+ Color-based criteria rave been used to identify quasars for ucarly LO vears (see the review by Burbidge 1967). aud the use of far-red filters to supplement blucr UDV oassbands was suggested as carly as 1967 (Braccesi 1967: Braccesi. Lyuds. Sandage 1968).," Color-based criteria have been used to identify quasars for nearly 40 years (see the review by Burbidge 1967), and the use of far-red filters to supplement bluer UBV passbands was suggested as early as 1967 (Braccesi 1967; Braccesi, Lynds, Sandage 1968)."
+ Quasars aro therefore extremely difficult to distinguish rou GRB afterelows using only broad band colors., Quasars are therefore extremely difficult to distinguish from GRB afterglows using only broad band colors.
+ The iain discriminant available is that quasars tend to be somewhat blucr than typical ὲtorelows., The main discriminant available is that quasars tend to be somewhat bluer than typical afterglows.
+ In practice. we expect confusion witli quasars to set the maxinuin area over which an unanibieuous afterglow identification cau be expected.," In practice, we expect confusion with quasars to set the maximum area over which an unambiguous afterglow identification can be expected."
+ This confusion cau be resolved with low resolution spectra. which identify most quasars unaunubieuouslv by their strong. broademission lines.," This confusion can be resolved with low resolution spectra, which identify most quasars unambiguously by their strong, broademission lines."
+ Sample color-color diagrams are shown for the pplaue in fgure l.. and the pplaue in figure 2..," Sample color-color diagrams are shown for the plane in figure \ref{mod_uvvi}, and the plane in figure \ref{mod_bvri}."
+ These lave been chosen as representative cases Wwrere the two-color method works., These have been chosen as representative cases where the two-color method works.
+ Additional coujuations of filters are also effective provided tha at least three filters are used. the bluest filter is at least as blue as the Johuson D haiC. aud the reddest is a least as red as Cousius I. Uufortuuatelv. the observationallv easiest filters BYR) do uot vield au adequate separation between afterglow aud stellar colors to expect reliable identification of afterelows. barring unusually Ligh photometric precision.," Additional combinations of filters are also effective provided that at least three filters are used, the bluest filter is at least as blue as the Johnson B band, and the reddest is at least as red as Cousins I. Unfortunately, the observationally easiest filters (BVR) do not yield an adequate separation between afterglow and stellar colors to expect reliable identification of afterglows, barring unusually high photometric precision."
+ Given a sufficient filter set. the afterglows are well separated from the stellar locus.," Given a sufficient filter set, the afterglows are well separated from the stellar locus."
+ Moreover. the reddening vector for Milkv Way dust runs essentially parallel to both the stellar sequence and the power law spectral sequence.," Moreover, the reddening vector for Milky Way dust runs essentially parallel to both the stellar sequence and the power law spectral sequence."
+ Thus. even though reddened afterglow spectra are no longer strict power laws. they remain distinct iu color-color space.," Thus, even though reddened afterglow spectra are no longer strict power laws, they remain distinct in color-color space."
+ Of ereater concern is dust in the CRB host ealaxv., Of greater concern is dust in the GRB host galaxy.
+ Because this dust is at ligh redshift. its 21715À ffcature can euter the observed wavelength range. changing the direction of the reddening vector in color-color space.," Because this dust is at high redshift, its $2175$ feature can enter the observed wavelength range, changing the direction of the reddening vector in color-color space."
+ For redshifts 2.2=23.3. this feature falls in the reddest flter of the color-color plot. thus making the loue-waveleneth color (c.e.. V-D slightly οι.," For redshifts $2.2 \la z \la 3.3$, this feature falls in the reddest filter of the color-color plot, thus making the long-wavelength color (e.g., V-I) slightly ."
+ A moderate amount of dust, A moderate amount of dust
+Changing integration variables to 7=gQ!+cot@. the augular integral becomes =2n.. aud we do not expect the total euerey in iucomipressive shiwaves to evolve.,"Changing integration variables to $\tau = q\Omega t + \cot\theta$, the angular integral becomes =, which is independent of time; hence = E_i (t = and we do not expect the total energy in incompressive shwaves to evolve."
+ Although this result may appear to depend iu detail on the assumption of isotropy. one can show that it really only depends on (Ελ=0)) being smooth near sin0=0. i.e. that there should uot be a conceutration of power in nearly radial wavevectors.," Although this result may appear to depend in detail on the assumption of isotropy, one can show that it really only depends on $\< E_{ki} (t =
+0)\>$ being smooth near $\sin \theta = 0$, i.e. that there should not be a concentration of power in nearly radial wavevectors."
+ This can be seeu from the following argument., This can be seen from the following argument.
+ [f we relax the assumptiou of isotropy. the angular iutegral becomes," If we relax the assumption of isotropy, the angular integral becomes."
+0)?.(33) For q0/x»1 the above iutegrande is sharply peaked in the narrow regionse around (ie.. sin02 0).," For $q\Omega t \gg 1$ the above integrand is sharply peaked in the narrow regions around $\tan \theta = -1/(q\Omega t) \ll 1$ (i.e., $\sin
+\theta \simeq 0$ )."
+ One can perform a Taylor-series expansion of (der(59.0); in these regions. and as long as (er(Ay.@)) itself is not sharply peaked it is well approximated as a constant.," One can perform a Taylor-series expansion of $\<
+\delta v_\perp^2 (k_0,\theta) \>$ in these regions, and as long as $\<
+\delta v_\perp^2(k_0,\theta) \>$ itself is not sharply peaked it is well approximated as a constant."
+ A modest relaxation of the assumption of isotropy. then. will result in an asyimnptotically constant value for the euergy integral.," A modest relaxation of the assumption of isotropy, then, will result in an asymptotically constant value for the energy integral."
+ Based upon this analysis. large amplification in an iucdividual siwave does uot iu itself argue lor a (rausition to turbulence due to transient growth.," Based upon this analysis, large amplification in an individual shwave does not in itself argue for a transition to turbulence due to transient growth."
+ Oue must also demoustrate that a “natural” set oL perturbations cau extract euergy from the backgrouud shear flow., One must also demonstrate that a “natural” set of perturbations can extract energy from the background shear flow.
+ Iu the case of the uustratified shearing sheet. the energy of a raucom set of incomipressive perturbations remalus constant. with time.," In the case of the unstratified shearing sheet, the energy of a random set of incompressive perturbations remains constant with time."
+ This is consistent with the results of Umurhau&οσον(2001).. who see asyinptotic decay in linear theory. because they work with a finite set of wavevectors. each of which must decay asyliptotically.," This is consistent with the results of \cite{ur04}, who see asymptotic decay in linear theory, because they work with a finite set of wavevectors, each of which must decay asymptotically."
+ Here we calculate the enereyOe evolution of the compressive shwaves for comparison purposes., Here we calculate the energy evolution of the compressive shwaves for comparison purposes.
+if stelar mass loss is recveled iuto new stars. so that e~l.,"if stellar mass loss is recycled into new stars, so that $a\sim 1$."
+ A inass-to-cherey efficiency of Wl has been used but it cau be 0.06 if the black tole is nof spline., A mass-to-energy efficiency of $0.1$ has been used but it can be $0.06$ if the black hole is not spinning.
+ or 0.12 if it )ocomies a maximiallv spinning. Wer. black hole.," or $0.42$ if it becomes a maximally spinning, Kerr, black hole."
+ A Level nore extreme possibility which defines au upper liuit on the efficiency relative to fio final (dead) black hole mass is to asstuue that the blac- hole was maximally sunius durius the accretion yhase and then spui down by. sav. the Dlaudford-Ziaes (L977) mechanisiu.," An even more extreme possibility which defines an upper limit on the efficiency relative to the final (dead) black hole mass is to assume that the black hole was maximally spinning during the accretion phase and then spun down by, say, the Blandford-Znajek (1977) mechanism."
+ The otal cucrey releascc relative to he final black hole mass allows for an order of maeutude πιοalntv 11a and tims £aev., The total energy released relative to the final black hole mass allows for an order of magnitude uncertainty in $\eta$ and thus $E_{AGN}$.
+ Of course a laugh value here. which maxiulses ονE.. overprediets tιο NRB inteusitv unCSS uost of the erowing phase of black ledes ds €‘ommptou thick.," Of course a high value here, which maximises $E_{AGN}/E_{\star}$, overpredicts the XRB intensity unless most of the growing phase of black holes is Compton thick."
+ It is also possible tha a significant fraction of f16 power fronan ACN is in he form of a wind aud not directly iu radiation., It is also possible that a significant fraction of the power from an AGN is in the form of a wind and not directly in radiation.
+ As ¢iscussed later. growing black holes may be both Comptoi thick and powerine wines.," As discussed later, growing black holes may be both Compton thick and powering winds."
+" If this is correct. then Eον/E, may be significantly higher than the estimate in the last seclon."," If this is correct, then $E_{AGN}/E_{\star}$ may be significantly higher than the estimate in the last section."
+ As outlined above. at least 85 per cent of accretion power is absorbed.," As outlined above, at least 85 per cent of accretion power is absorbed."
+ Since about ten per cel 1511 quasars whici show very little absorption. this mieaus tha lost lines of sight ou of the remaining objects are highly absorbed.," Since about ten per cent is in quasars which show very little absorption, this means that most lines of sight out of the remaining objects are highly absorbed."
+ This is difficult ‘OL the standard obscuring torus model. which could absorb perhaps oue half ο two thirds of al sigoyit lines.," This is difficult for the standard obscuring torus model, which could absorb perhaps one half to two thirds of all sight lines."
+" Even hen it is unclear what inflates the torus. wuch Is supposecl Τε) be cold. aux molecular,"," Even then it is unclear what inflates the torus, which is supposed to be cold and molecular."
+ Dissipation in in a system of orbiting clotds should cause i o flaten iuto a disc. with lowcovering factor.," Dissipation in in a system of orbiting clouds should cause it to flatten into a disc, with lowcovering factor."
+ Encrev unst be οςntiuuousVv injected into anv cold absorbing cloud svsCll to keep it 1iflated and soszv covering., Energy must be continuously injected into any cold absorbing cloud system to keep it inflated and so sky covering.
+ One plausible sohtion is that a gas-rkles ar cluster sirenids the lack hole aud it is the massive stars (winds and superuovae) which supply the energye (Fabia ret al 1998)., One plausible solution is that a gas-rich star cluster surrounds the black hole and it is the massive stars (winds and supernovae) which supply the energy (Fabian et al 1998).
+ The surrouding starburst cant iereby Obscure he acfive nuc‘leus, The surrounding starburst can thereby obscure the active nucleus.
+ The stilirst shotld chhaunce the metallicity of the absorbing eas., The starburst should enhance the metallicity of the absorbing gas.
+ This makes a oejiven nass of easone nore efficient at absorbing N-ravs iux indeed increases the effect O absorption before Compton down-scattering comes iito play., This makes a given mass of gas more efficient at absorbing X-rays and indeed increases the effect of absorption before Compton down-scattering comes into play.
+ This is iuportaut 1n openiue up the yaraineter space for modeL-fitting of 10 NRB spectrum AWVilman Fabian 1999)., This is important in opening up the parameter space for model-fitting of the XRB spectrum (Wilman Fabian 1999).
+ Fuclling of the wecleus is au old problem (see e.g. Shostman et al 1990)., Fuelling of the nucleus is an old problem (see e.g. Shlosman et al 1990).
+ Although there nav be lots of eas around the nucleus. augular omentum nav prevent it from rapidly accreΠιο to the centre.," Although there may be lots of gas around the nucleus, angular momentum may prevent it from rapidly accreting to the centre."
+ In this respect. a hot phase in the surrounding medi may beimiportanut. with Boudi accretion from fis phase being the dominant mechanisi (see e.c. Nulseu Fabian 1999).," In this respect, a hot phase in the surrounding medium may be important, with Bondi accretion from this phase being the dominant mechanism (see e.g. Nulsen Fabian 1999)."
+ Aneular momenta may be transported outward bv turbulence within such a hot phase. so allowing rapid accretion to proceed.," Angular momentum may be transported outward by turbulence within such a hot phase, so allowing rapid accretion to proceed."
+"the other hand, a small vsin does not prevent mapping using Zeeman-Doppler imaging as described in Sect.","the other hand, a small $v\sin i$ does not prevent mapping using Zeeman-Doppler imaging as described in Sect."
+ 3. (, 3. (
+e.g. Petit et al.,e.g. Petit et al.
+ 2008)., 2008).
+" Figure 2 shows that the seasonal variations of B; and the activity indicators generally vary coherently with photometric phase, and that there is seasonal scatter greater than the measurement error bars."," Figure 2 shows that the seasonal variations of $B_{\ell}$ and the activity indicators generally vary coherently with photometric phase, and that there is seasonal scatter greater than the measurement error bars."
+ In addition there are some outstanding deviations from the mean variation., In addition there are some outstanding deviations from the mean variation.
+" As far as the Bg measurements are concerned, the observations of 20 September 2007 (the first one, which is our “discovery observation""), as well as those obtained on 21 September and 13 October 2010 present outstanding deviations from the general phase variations."," As far as the $B_{\ell}$ measurements are concerned, the observations of 20 September 2007 (the first one, which is our “discovery observation”), as well as those obtained on 21 September and 13 October 2010 present outstanding deviations from the general phase variations."
+" On 20 September 2007, the S-index and other activity indicators do not deviate from the general phase variation."," On 20 September 2007, the $S$ -index and other activity indicators do not deviate from the general phase variation."
+" On the other hand, for the 2 first autumn 2010 observations, both By and activity indicators are stronger than expected at their respective phases."," On the other hand, for the 2 first autumn 2010 observations, both $B_{\ell}$ and activity indicators are stronger than expected at their respective phases."
+" Remarkably, in the 2010-2011 season, the modulation of activity indicators disappears: this season is highlighted in red in Fig."," Remarkably, in the 2010-2011 season, the modulation of activity indicators disappears: this season is highlighted in red in Fig."
+ 2., 2.
+" While B; values approach the typical level and modulation at the end of 2010, the intensity of the activity indicators remains high."," While $B_{\ell}$ values approach the typical level and modulation at the end of 2010, the intensity of the activity indicators remains high."
+" Also, on some occasions, in particular 28 September 2009, the activity indicators are surprisingly high, while the Be value is quite normal with respect to the phase."," Also, on some occasions, in particular 28 September 2009, the activity indicators are surprisingly high, while the $B_{\ell}$ value is quite normal with respect to the phase."
+" Looking closely at the data, we find that we may have a problem of normalization of the continuum for this date, which can affect the measurements of activity indicators, but not B, (as the Stokes V measurement is a differential measurement)."," Looking closely at the data, we find that we may have a problem of normalization of the continuum for this date, which can affect the measurements of activity indicators, but not $B_{\ell}$ (as the Stokes $V$ measurement is a differential measurement)."
+" 'To further investigate this potential background problem, we have computed the Ca H emission index used by Morgenthaler et al. ("," To further investigate this potential background problem, we have computed the Ca H emission index used by Morgenthaler et al. ("
+"2011, 2012).","2011, 2012)."
+ This method makes use of a synthetic spectrum from the POLLUX database (Palacios et al., This method makes use of a synthetic spectrum from the POLLUX database (Palacios et al.
+ 2010) to normalise the continuum., 2010) to normalise the continuum.
+" It was found to be very effective in the case of solar-type dwarfs, enabling those authors to reach internal errors of about 0.001 for chromospheric emission index measurements (Morgenthaler et al."," It was found to be very effective in the case of solar-type dwarfs, enabling those authors to reach internal errors of about 0.001 for chromospheric emission index measurements (Morgenthaler et al."
+" 2011, 2012)."," 2011, 2012)."
+" However, in the case of EK Eri this method did not allow us to reduce the deviations, and we present our S-index measurements in Table 1 and Fig."," However, in the case of EK Eri this method did not allow us to reduce the deviations, and we present our $S$ -index measurements in Table 1 and Fig."
+ 2 as they are., 2 as they are.
+" Of peculiar interest are the isolated enhancements of Bz, activity indicators or both."," Of peculiar interest are the isolated enhancements of $B_{\ell}$, activity indicators or both."
+ They may be associated with flares as observed in active giants (e.g. Konstantinova-Antova et al., They may be associated with flares as observed in active giants (e.g. Konstantinova-Antova et al.
+" 2000, 2005)."," 2000, 2005)."
+ In this case it would be the first time that simultaneous observations of magnetic field and activity indicators have been performed during flares in the stellar context., In this case it would be the first time that simultaneous observations of magnetic field and activity indicators have been performed during flares in the stellar context.
+" In the solar case, a simultaneous increase of magnetic field and Ho has been observed (Lozitsky et al."," In the solar case, a simultaneous increase of magnetic field and $\alpha$ has been observed (Lozitsky et al."
+ 2000)., 2000).
+ The outlying magnetic observations of September 2007 and autumn 2010 were found to increase dramatically the reduced χ of our model fitting in Sect., The outlying magnetic observations of September 2007 and autumn 2010 were found to increase dramatically the reduced $\chi^2$ of our model fitting in Sect.
+ 3., 3.
+ We will use only data between November 2007 and March 2010 (spanning about 3 rotations) for forthcoming magnetic analysis., We will use only data between November 2007 and March 2010 (spanning about 3 rotations) for forthcoming magnetic analysis.
+" In order to model our series of Stokes V line profiles and to reconstruct the surface magnetic geometry of the star, we have used the Zeeman-Doppler imaging inversion method (ZDI; Donati and Brown 1997)."," In order to model our series of Stokes $V$ line profiles and to reconstruct the surface magnetic geometry of the star, we have used the Zeeman-Doppler imaging inversion method (ZDI; Donati and Brown 1997)."
+ The version of the code employed is that described by Donati et al. (, The version of the code employed is that described by Donati et al. (
+2006b): the surface magnetic field is projected onto a spherical harmonics frame and the magnetic field is resolved into poloidal and toroidal components.,2006b): the surface magnetic field is projected onto a spherical harmonics frame and the magnetic field is resolved into poloidal and toroidal components.
+" Because of the very slow rotation of EK Eri, its vsin has been impossible to determine up to now due to the degeneracy between vsini and macroturbulence (DBSS10)."," Because of the very slow rotation of EK Eri, its $v\sin i$ has been impossible to determine up to now due to the degeneracy between $vsini$ and macroturbulence (DBSS10)."
+ These authors consider that 0.8 km s! is a safe upper limit for vsini., These authors consider that 0.8 km $^{-1}$ is a safe upper limit for $v\sin i$.
+ In our models we used vsini = 0.5 km s! and found that varying this value did not significantly change our results., In our models we used $v\sin i$ = 0.5 km $^{-1}$ and found that varying this value did not significantly change our results.
+ We used a linear limb darkening coefficient equal to 0.75., We used a linear limb darkening coefficient equal to 0.75.
+ We limited the spherical harmonics expansion to £<3 since increasing this threshold do not significantly change the results., We limited the spherical harmonics expansion to $\ell \le 3$ since increasing this threshold do not significantly change the results.
+ Knowing the rotational period of EK Eri is essential to be able to perform ZDI., Knowing the rotational period of EK Eri is essential to be able to perform ZDI.
+ Photometric variations of EK Eri have been now monitored during more than 30 years and a period of 308.8 + 2.5 days has been adopted (DBSS10) as well as an ephemeris., Photometric variations of EK Eri have been now monitored during more than 30 years and a period of 308.8 $\pm$ 2.5 days has been adopted (DBSS10) as well as an ephemeris.
+ This period may, This period may
+The recent neutrino oscillation data have greatly. enriched. our knowledge of the flavor structure of leptons.,The recent neutrino oscillation data have greatly enriched our knowledge of the flavor structure of leptons.
+ After the discovery of neutrino oscillations. there has," After the discovery of neutrino oscillations, there has"
+"(2004) measure an integrated. fluxdensitv of == 290417Jv [for the ""e pair. named Ηλ JO40343 in the LUPASS Bright Galaxy Catalog(see Table 1).","(2004) measure an integrated flux density of = $259\pm17$ for the galaxy pair, named HIPASS J0403–43 in the HIPASS Bright Galaxy Catalog (see Table 1)."
+ The detected cemission is centered. on NGC 1519 ancl significantly extended with respect to the Parkes gridded beum of 15/55., The detected emission is centered on NGC 1512 and significantly extended with respect to the Parkes gridded beam of 5.
+ llawarden et al. (, Hawarden et al. (
+1979) measured —— 232+20 ((same as Reif et al.,1979) measured = $232\pm20$ (same as Reif et al.
+ 1982). slightly lower than the LILPASS value.," 1982), slightly lower than the HIPASS value."
+ Llere we present high-resolution ATCA Iline and 20-em radio continuum data of the galaxy pair NGC 1512/1510 as well as complimentary GALEN CY-. SINGG Lla-- and Spitzer mid-infrared. images.," Here we present high-resolution ATCA line and 20-cm radio continuum data of the galaxy pair NGC 1512/1510 as well as complimentary GALEX $UV$ -, SINGG - and Spitzer mid-infrared images."
+ The paper is organised. as follows: in Section 2 we summarise the observations and data reduction: in Section. 3 we present the lline ancl the 20-cm radio continuum results. including ο discovery of twogalary candidates.," The paper is organised as follows: in Section 2 we summarise the observations and data reduction; in Section 3 we present the line and the 20-cm radio continuum results, including our discovery of two candidates."
+ The discussion. in Section 4 exploits the available wavelength data sets. comparing the eeasoo density with the properties of star-forming regions out to radii of SO kpc.," The discussion in Section 4 exploits the available multi-wavelength data sets, comparing the gas density with the properties of star-forming regions out to radii of 80 kpc."
+ Section 5 contains our conclusions and Section6 a brief outlook towards ssurvevs with the Australian SIA Pathfinder CASILAD)., Section 5 contains our conclusions and Section 6 a brief outlook towards surveys with the Australian SKA Pathfinder (ASKAP).
+ {line and Ἄθ-ομι racio continuum observations ol he galaxy pair NGC. 1512/1510 were. obtained with he Australia Telescope Compact Array (ATCA) using multiple configurations and four (overlapping) pointings., line and 20-cm radio continuum observations of the galaxy pair NGC 1512/1510 were obtained with the Australia Telescope Compact Array (ATCA) using multiple configurations and four (overlapping) pointings.
+ The observing details are given in Table 2., The observing details are given in Table 2.
+ The first. frequency band. (IE1) was centered on 1415 MlIz with a bandwicth of S Mllz. divided into 512 channels.," The first frequency band (IF1) was centered on 1415 MHz with a bandwidth of 8 MHz, divided into 512 channels."
+ ‘This gives a channel width of aand a velocity resolution of, This gives a channel width of and a velocity resolution of.
+ The ATCA primary beam is 33/66 at 1415 ΜΗΝ., The ATCA primary beam is 6 at 1415 MHz.
+ second. frequency. band (1E2) was centered on 1384 MILTIzMetsthe(em) with a bandwidth of 128 Mllz civided into 32 , The second frequency band (IF2) was centered on 1384 MHz (20-cm) with a bandwidth of 128 MHz divided into 32 channels.
+The APCA is a radio interferometer consisting ol six 22-m clishes. creating 15 baselines in a single configuration. equipped with seven receiver systems covering wavelengths from 3-mum to 20-cm.," The ATCA is a radio interferometer consisting of six 22-m dishes, creating 15 baselines in a single configuration, equipped with seven receiver systems covering wavelengths from 3-mm to 20-cm."
+ While five antennas (CAO01 to €A05) are movable along a 3-km long cast-west track (and a 214-2 long north-south spur. allowing us to create hybrid arrays). one antenna (CA06) is fixed at a distance of 3-km from the end of the track.," While five antennas (CA01 to CA05) are movable along a 3-km long east-west track (and a 214-m long north-south spur, allowing us to create hybrid arrays), one antenna (CA06) is fixed at a distance of 3-km from the end of the track."
+ By combining data from several array configuration (see Table 2) we achieve excellent. ne-coverage eencrated by over 100 basclines ranging from 30-m to 6-ko., By combining data from several array configuration (see Table 2) we achieve excellent $uv$ -coverage generated by over 100 baselines ranging from 30-m to 6-km.
+ Using Fourier transformation. this allows us to make data cubes and images at a laree range of angular resolutions (up to aat 20-cm) by choosing cillerent weights for short. medium and long baselines which in turn are sensitive to cilferent structure scales.," Using Fourier transformation, this allows us to make data cubes and images at a large range of angular resolutions (up to at 20-cm) by choosing different weights for short, medium and long baselines which in turn are sensitive to different structure scales."
+ The weighting of the data allects not only the resolution. but also the rms noise and sensitivity to clilfuse emission.," The weighting of the data affects not only the resolution, but also the rms noise and sensitivity to diffuse emission."
+ Data reduction was carried out. with the £ware package (Sault. Teuben Wright 1995) using μαancdard: procedures.," Data reduction was carried out with the software package (Sault, Teuben Wright 1995) using standard procedures."
+ After calibration the IET data were Spit into a narrow band 20-em racio continuum and an Iline clata set using a first order fit to the line-free channels., After calibration the IF1 data were split into a narrow band 20-cm radio continuum and an line data set using a first order fit to the line-free channels.
+" ccubes were mace using ""natural, (na) and ""robust? (r=0) weighting of the wecata in the velocity range covered. by the comission using steps of", cubes were made using `natural' (na) and `robust' (r=0) weighting of the -data in the velocity range covered by the emission using steps of.
+ The longest. baselines to the distant antenna six (CA06) were excluded: when making the low-resolution cubes., The longest baselines to the distant antenna six (CA06) were excluded when making the low-resolution cubes.
+" Broad-bancl 20-cmi racio continuum images were made using “robust” (r—0) and ""uniform weighting of the IE2 sedata.", Broad-band 20-cm radio continuum images were made using `robust' (r=0) and `uniform' weighting of the IF2 -data.
+ The data were analvsed. usingMIRIAD. apart from the rotation curve Lit which was obtained using the software package (van der Hulst et αἱ.," The data were analysed using, apart from the rotation curve fit which was obtained using the software package (van der Hulst et al."
+ 1992)., 1992).
+ The NGC 1512/1510 galaxy pair is an impressive svstenm., The NGC 1512/1510 galaxy pair is an impressive system.
+" Our ALCA mnmiosaic (see Flies,", Our ATCA mosaic (see Figs.
+ 4) shows a very extended. gas distribution. spanning a diameter of 40 (for 110 kpe).," 2–4) shows a very extended gas distribution, spanning a diameter of $\sim$ (or 110 kpc)."
+ Two prominent spiral arms. which appear to," Two prominent spiral arms, which appear to"
+et al.,et al.
+ 2006b)., 2006b).
+ Within this sample there are 2140 EROs with shotometric redshifts in the range OS«zld and z/<25., Within this sample there are 2140 EROs with photometric redshifts in the range $0.8<z<1.4$ and $z^{\prime}\leq25$.
+ As expected. none of these EROs display a statistically acceptable shotometric redshift solution at 2>5.," As expected, none of these EROs display a statistically acceptable photometric redshift solution at $z\geq5$."
+ The fundamental reason for 1s is that none of them have the blue SED slope long-ward of the z/ band displayed by our final =>5 LBG candidates., The fundamental reason for this is that none of them have the blue SED slope long-ward of the $z^{\prime}-$ band displayed by our final $z\geq5$ LBG candidates.
+" Perhaps 1e most convincing evidence for this comes from the SED of the stacked photometry of all nine 2=5 LBG candidates (Figs 2 3). which has a colour of z—LA,=0.73+ 0.25."," Perhaps the most convincing evidence for this comes from the SED of the stacked photometry of all nine $z\geq5$ LBG candidates (Figs 2 3), which has a colour of $z^{\prime}-K=0.73\pm0.25$ ."
+ For comparison. re average οA colour of =~1 EROs in the UDS field is So/dy=2.3240.32 (where the quoted uncertainty is the standard deviation).," For comparison, the average $z^{\prime}-K$ colour of $z\simeq1$ EROs in the UDS field is $z^{\prime}-K=2.32\pm0.32$ (where the quoted uncertainty is the standard deviation)."
+ High-redshift quasars in the redshift range 5.0«z6.5 have predicted //.— z'and z/J colours which are comparable with those of our tinal nine LBG candidates (Willott et al., High-redshift quasars in the redshift range $5.0<z<6.5$ have predicted $i^{\prime}-z^{\prime}$ and $z^{\prime}-J$ colours which are comparable with those of our final nine LBG candidates (Willott et al.
+ 2005)., 2005).
+ However. none of the tinal candidates are detected in the deep 0.212 keV band X-ray imaging of the SXDS with XMM-Newton.," However, none of the final candidates are detected in the deep $0.2-12$ keV band X-ray imaging of the SXDS with XMM-Newton."
+ The sensitivity of the XMM-Newton SXDS observations varies somewhat over the SXDS field. but the typical X-ray flux upper limits imply X-ray luminosity limits of &(5—10)«107! erg + (rest-frame energy band ~1TO keV).," The sensitivity of the XMM-Newton SXDS observations varies somewhat over the SXDS field, but the typical X-ray flux upper limits imply X-ray luminosity limits of $\simeq
+(5-10)\times10^{44}$ erg $^{-1}$ (rest-frame energy band $\sim1-70$ keV)."
+ Moreover. based on the latest constraints on the QSO luminosity function at 2—6 (Willott et al.," Moreover, based on the latest constraints on the QSO luminosity function at $z\simeq6$ (Willott et al."
+ 2005). contamination of our sample from high-redshift QSOs is expected to be minimal (conservatively 1 QSO is anticipated with z<25 over the SXDS/UDS field).," 2005), contamination of our sample from high-redshift QSOs is expected to be minimal (conservatively $\leq1$ QSO is anticipated with $z^{\prime}\leq25$ over the SXDS/UDS field)."
+ Finally. given that we have deliberately selected the brightest objects of their type at 2—5. it is important to at least consider the possibility that some of the tinal LBG candidates may be lensed.," Finally, given that we have deliberately selected the brightest objects of their type at $z\geq5$, it is important to at least consider the possibility that some of the final LBG candidates may be lensed."
+ The first scenario to consider is the prospect of moderate amplification due to a lensing geometry where the LBG candidates lie outside the Einstein radii of lensing objects along the line-of-sight., The first scenario to consider is the prospect of moderate amplification due to a lensing geometry where the LBG candidates lie outside the Einstein radii of lensing objects along the line-of-sight.
+" Tf we model the potential lens as a single isothermal sphere. then the amplification is : //=ih0,7 where 6p is the Einstein radius @2=tifoy2e- Gy is the velocity dispersion of the lensing galaxy. efQs Dsis the angular diameter distance of the source and {δες is the angular diameter distance of the source from the lens."," If we model the potential lens as a single isothermal sphere, then the amplification is : $\mu=\frac{\theta}{\theta-\theta_{E}}$, where $\theta_{E}$ is the Einstein radius $\theta_{E}=4\pi(\frac{\sigma_{V}}{c})^2\frac{D_{LS}}{D_{S}}$, $\sigma_{V}$ is the velocity dispersion of the lensing galaxy, $D_{S}$ is the angular diameter distance of the source and $D_{LS}$ is the angular diameter distance of the source from the lens."
+" Four of the tinal nine LBG candidates have nearby companions within 2.5""of the line-of-sight. which could potentially produce amplification via lensing."," Four of the final nine LBG candidates have nearby companions within $2.5$ of the line-of-sight, which could potentially produce amplification via lensing."
+ To investigate this we have estimated the velocity dispersionsand consequent amplification for these objects. using their photometric redshifts and the AJ;σι; correlation at 2οO from Bernardi et al. (," To investigate this we have estimated the velocity dispersionsand consequent amplification for these objects, using their photometric redshifts and the $M_{i}-\sigma_{V}$ correlation at $z\simeq0$ from Bernardi et al. ("
+2003).,2003).
+ This calculation reveals that none of the companion objects are compatible with having a25.100 km +. and that any amplification from lensing is therefore small (ie. jr— 1.1).," This calculation reveals that none of the companion objects are compatible with having $\sigma_{V}\geq100$ km $^{-1}$, and that any amplification from lensing is therefore small (i.e. $\mu\leq1.1$ )."
+ The second scenario to consider is the prospect of high amplification due to extremely close alignment along the line-of-sight between the LBG candidates and potential lensing objects in the foreground., The second scenario to consider is the prospect of high amplification due to extremely close alignment along the line-of-sight between the LBG candidates and potential lensing objects in the foreground.
+ Although difficult to rule-out completely. the remarkably clean nature of the LBG candidates’ optical photometry suggests thatit is quite unlikely.," Although difficult to rule-out completely, the remarkably clean nature of the LBG candidates' optical photometry suggests thatit is quite unlikely."
+" For the 25 LBG candidates the largest cross-section for lensing occurs at >&0.5,", For the $z\geq5$ LBG candidates the largest cross-section for lensing occurs at $z\simeq0.8$.
+" Consequently. if we consider an £L"" elliptical at 2=0.5 as the potential lens. we can predict that it would have apparent magnitudes of ἐν&19.9. Vox25.5 and D=26.6 (assuming Aj=—22.9: Pozzetti et al."," Consequently, if we consider an $L^{\star}$ elliptical at $z=0.8$ as the potential lens, we can predict that it would have apparent magnitudes of $K\simeq19.9$, $V\simeq25.5$ and $B\simeq26.6$ (assuming $M^{\star}_{K}=-22.9$; Pozzetti et al."
+ 2003)., 2003).
+ These magnitudes are completely incompatiblewith those of the stack of the final nine LBG candidates. which is undetected to 2a limits ofB=20.5 and V= 25.7. and has ads band magnitude of AV=23.86+ 0.23.," These magnitudes are completely incompatiblewith those of the stack of the final nine LBG candidates, which is undetected to $2\sigma$ limits of$B=29.5$ and $V=28.7$ , and has a $K-$ band magnitude of $K=23.86\pm0.23$ ."
+ Consequently. the average luminosity of any potential lensing objects would have to," Consequently, the average luminosity of any potential lensing objects would have to"
+to be 10 un. since the mask contours can be customized to the as-built focal plane mechanical interface.,"to be $\sim 10\,\mu$ m, since the mask contours can be customized to the as-built focal plane mechanical interface."
+ Thus. 17prim are allowed for component (2) ancl The mask surfaces shall have the lowest possible reflectivity at the operating wavelengths (370 to 1850 nia).," Thus, $\sim 17\,\mu$ m are allowed for component (2) and The mask surfaces shall have the lowest possible reflectivity at the operating wavelengths (370 to 1850 nm)."
+ To miuunize observing overlreads. aud to avoid people inside the domes during the night. MOS iustments at the VLT must be able to change the mask configuration through remote control: iu our case. both VIMOS aud NIRMOS cau have 15 dillerent mask sets at any time.," To minimize observing overheads, and to avoid people inside the domes during the night, MOS instruments at the VLT must be able to change the mask configuration through remote control: in our case, both VIMOS and NIRMOS can have 15 different mask sets at any time."
+ For each iustrurent and for each quadrant there is a demountable Lustrument Cabinet (IC) that must ye prepared by the MAL operator by inserting the requested masks as much in advance of the spectroscopxc observations as possible., For each instrument and for each quadrant there is a demountable Instrument Cabinet (IC) that must be prepared by the MMU operator by inserting the requested masks as much in advance of the spectroscopic observations as possible.
+ These requirements imply that a large number of masks is sept availalje at any time. (rom which to elioose the oues to be inserted iu the ICs for the following üeht(5) of observation.," These requirements imply that a large number of masks is kept available at any time, from which to choose the ones to be inserted in the ICs for the following night(s) of observation."
+ Storage for 100 mask sets (100 masks) has » be provided., Storage for 100 mask sets (400 masks) has to be provided.
+ The masks must hus be uniquely identilied. must be traceable aud their position in each IC must be certain aud nade known to instrunent control software.," The masks must thus be uniquely identified, must be traceable and their position in each IC must be certain and made known to instrument control software."
+ A mask handliug system. providing hardware aud software tools to reach the goal of having the right mask in the focal plane at the right time. must ye In the initial concept (1997). the MMAL was intended as a milling machine which would cut the slits in a O.1 ium thin brass sheet. supported by a LO min thick aluminum frame.," A mask handling system, providing hardware and software tools to reach the goal of having the right mask in the focal plane at the right time, must be In the initial concept (1997), the MMM was intended as a milling machine which would cut the slits in a 0.1 mm thin brass sheet, supported by a 10 mm thick aluminum frame."
+ We assembled a simall milliug machine aud we sroved that the rougliness aud speed requirements were fulfilled., We assembled a small milling machine and we proved that the roughness and speed requirements were fulfilled.
+" The minimum obtainable slit wih was 300 jin. This solution was discarded because the ICs were too large and the accuracy in the slit positioning was hampered by the tlieri expausion of brass. eiven the j»ossible temperature «ifIereuces between the time a mask was imauufactured anc usec in the iust""nent or the temperaure variations during the observations."," The minimum obtainable slit width was $300\,\mu$ m. This solution was discarded because the ICs were too large and the accuracy in the slit positioning was hampered by the thermal expansion of brass, given the possible temperature differences between the time a mask was manufactured and used in the instrument or the temperature variations during the observations."
+ Furthermore the lifetime of the cutting tools was short cle to breaking aud to wear (that caused a slit width variation)., Furthermore the lifetime of the cutting tools was short due to breaking and to wear (that caused a slit width variation).
+ Subsecquenu developments (1998) were aimed at iuininising the sources of errors. by makine use of altunintin and thicker (but sill <0.3 mun). frameless masks. which had the advantage of reclucing tje size aud the weight ol the ICs.," Subsequent developments (1998) were aimed at minimising the sources of errors, by making use of aluminum and thicker (but still $<\,0.3$ mm), frameless masks, which had the advantage of reducing the size and the weight of the ICs."
+ The next natural step was the use of a immaterial with a low therual expansion coellicieu: we tested carbon fiber. kevlar. eraphite and invar. but it was," The next natural step was the use of a material with a low thermal expansion coefficient: we tested carbon fiber, kevlar, graphite and invar, but it was"
+X-ray eniission.,X-ray emission.
+" For p=2.2. the timescale Z5; is givenbv: where [ is the bulk lorentz factor. 5,,5, is the minimum lorentz factor of the accelerated electron by turbulence and » is the number density of shocked medium."," For $p=2.2$, the timescale $T_d$ is givenby: where $\Gamma$ is the bulk lorentz factor, $\gamma_{min}$ is the minimum lorentz factor of the accelerated electron by turbulence and $n$ is the number density of shocked medium."
+ The dissipation timescale of the turbulent energy is coincicent with the observational timescale of shallow decay. phase., The dissipation timescale of the turbulent energy is coincident with the observational timescale of shallow decay phase.
+ Therefore. in the shocked region with turbulence. (he evolution ofmicrophysical parameters « τη aud 2; can be derived rom the coefficients of Fokker-Planck equation.," Therefore, in the shocked region with turbulence, the evolution of microphysical parameters $\varepsilon_e$, $\varepsilon_B$ and $\varepsilon_t$ can be derived from the coefficients of Fokker-Planck equation."
+" From equation (7) ancl (8). we obtain: and Furthermore. in this turbulent region.ex we assume that these microphysical parameters have the forms. ofn z,x(. 2gLxll and ο,x(7, "," From equation (7) and (8), we obtain: and Furthermore, in this turbulent region, we assume that these microphysical parameters have the forms of $\varepsilon_e\propto
+t^a$, $\varepsilon_B\propto t^b$ and $\varepsilon_t\propto t^d$."
+We- insert. them into. equation. (11)x and (12)EV and put the constraints on a. b and d by (wo algebra equations: and We choose the turbulent spectrum of Ixraichnan with the index q=4/3.," We insert them into equation (11) and (12) and put the constraints on a, b and d by two algebra equations: and We choose the turbulent spectrum of Kraichnan with the index $q=4/3$."
+ All of the possible values of a. b and d are shown in Figure 1. 2 and 3.," All of the possible values of $a$, $b$ and $d$ are shown in Figure 1, 2 and 3."
+ From the standard afterglow model. the fluxes in the early X-ray band. were written by Sari.Piran&Naravan(1993).," From the standard afterglow model, the fluxes in the early X-ray band were written by \citet{sari98}."
+. There are two different limits: acliabatic and radiative case., There are two different limits: adiabatic and radiative case.
+ For explanation of (he emission in shallow decay phase. we obtain the early X-ray light curve under the adiabatic case as: As an example. we adopt the index g=4/23 and p=2.2.," For explanation of the emission in shallow decay phase, we obtain the early X-ray light curve under the adiabatic case as: As an example, we adopt the index $q=4/3$ and $p=2.2$."
+ Four evolutionary [Iuxes are achieved in the table 1.. they are corresponding to the minimum and maxinunm values of e and 5.," Four evolutionary fluxes are achieved in the table \ref{tbl-1}, they are corresponding to the minimum and maximum values of $a$ and $b$ ."
+ While for the radiative case. (he X-ray flux is: Table 2. lists the possible fixes evolved with the time.," While for the radiative case, the X-ray flux is: Table \ref{tbl-2} lists the possible fluxes evolved with the time."
+The evolution of these vortices is ultimately related (o the dvnamies of convective motions in the domain.,The evolution of these vortices is ultimately related to the dynamics of convective motions in the domain.
+ The convective flows may sometime collect swirls in a local area. then merge and destrov them.," The convective flows may sometime collect swirls in a local area, then merge and destroy them."
+ Such vortical structures in simulations were first described by Stein(1998)., Such vortical structures in simulations were first described by \cite{stein1998}.
+. They showed (hat stronger vortices usually correlate with downllows. ancl (his is also found in our results.," They showed that stronger vortices usually correlate with downflows, and this is also found in our results."
+ From time to time. convection creates pretty big whirlpools. as the one indicated bv square in Fig. l..," From time to time, convection creates pretty big whirlpools, as the one indicated by square in Fig. \ref{noMHDswirls},"
+ which can swallow up other smaller swirls around them., which can swallow up other smaller swirls around them.
+ The big swirls are usually easy (o see also in (he surface temperature and intensity. variations., The big swirls are usually easy to see also in the surface temperature and intensity variations.
+ The detailed structure of a large whirlpool is shown in Figs lec-f The whirlpool structure is characterized bv: 1) formation of “arms” of higher density (hat correlates with lower temperature: 2) a pronounced vorlical structure of the velocity Mow: 3) increased magnitude of the horizontal velocity upto T. 9 km/s: 4) a sharply decreased density in the central core of the vortex. and a sliehtly higher temperature than in (he surrounding.," The detailed structure of a large whirlpool is shown in Figs \ref{noMHDswirls}c c-f. The whirlpool structure is characterized by: 1) formation of ”arms” of higher density that correlates with lower temperature; 2) a pronounced vortical structure of the velocity flow; 3) increased magnitude of the horizontal velocity up to 7 – 9 km/s; 4) a sharply decreased density in the central core of the vortex, and a slightly higher temperature than in the surrounding."
+ The tvpical depth of large swirls is about 100—200 kan., The typical depth of large swirls is about 100 – 200 km.
+ Inside the whirlpool shown in Fig., Inside the whirlpool shown in Fig.
+ lee. we can see a higher temperature (ube-like structure. but it is unstable (in comparison with the whirlpool lifetime). and can be destroved during the swirling motions.," \ref{noMHDswirls}e e, we can see a higher temperature tube-like structure, but it is unstable (in comparison with the whirlpool lifetime), and can be destroyed during the swirling motions."
+ The vorücal motions in the solar granulation have been detected in high-resolution observations (e.g.Pótzi&Brandt2005).. and the observational results generally agree with the simulations.," The vortical motions in the solar granulation have been detected in high-resolution observations \citep[e.g.][]{potzi2005}, and the observational results generally agree with the simulations."
+ In. particular. recent observations of a quiet region. near the solar disk center detected magnetic bright points following a logarithmic spiral trajectory around intergranular points and engulled bv a downdralt (Bonetetal.2008).," In particular, recent observations of a quiet region, near the solar disk center detected magnetic bright points following a logarithmic spiral trajectory around intergranular points and engulfed by a downdraft \citep{bonet08}."
+. The observations were interpreted. as vortical flows (hat affects the bright point motions., The observations were interpreted as vortical flows that affects the bright point motions.
+ These whirlpools have the size <0.5 Mm ancl the lifetime of about 5 min. without preferred. direction of rotation.," These whirlpools have the size $\lesssim 0.5$ Mm and the lifetime of about 5 min, without preferred direction of rotation."
+ The distribution of vortices studied from the ground (Wangetal.1995:Potzi&Brandt2005.2007:Bonetetal.2008) and space observations (Attieetal.2009) shows strong prelerences to concentration in regions of downflows. parliculary at the mesogranular scale (Pótzi&Branclt2005.2007).," The distribution of vortices studied from the ground \citep{wang1995,potzi2005,potzi2007,bonet08}
+ and space observations \citep{attie2009} shows strong preferences to concentration in regions of downflows, particulary at the mesogranular scale \citep{potzi2005,potzi2007}."
+. Our simulations for the domain of the horizontal size of 12.8 Mm also show a tendency of concentration of vortices on a mesogranular scale., Our simulations for the domain of the horizontal size of 12.8 Mm also show a tendency of concentration of vortices on a mesogranular scale.
+ We plan {ο discuss this effect in a separate paper., We plan to discuss this effect in a separate paper.
+ Here we focus on the links between the whirlpools and magnetic structure formation., Here we focus on the links between the whirlpools and magnetic structure formation.
+" To investigate (he process of magnetic field structuring in the turbulent. convective plasma we mace a series of simulations for (he initial vertical uniform magnetic field. τμ, varving from 1 to 100 G. different computational grids and domain sizes."," To investigate the process of magnetic field structuring in the turbulent convective plasma we made a series of simulations for the initial vertical uniform magnetic field, $Bz_0$, varying from 1 to 100 G, different computational grids and domain sizes."
+ Qualitatively the, Qualitatively the
+"the followings: (1) For the two photospheric lines, the spatial distribution of penumbral Evershed flow in 1; Dopplergrams","the followings: (1) For the two photospheric lines, the spatial distribution of penumbral Evershed flow in $\nu_i$ Dopplergrams"
+NelF. Relation (1)) aakes it possible to interpret disk fracinentation occurmiug at teuο as the inability of the disk to sustain gravitational stress at a21 (Rice 2005).,"c_s^2/F. Relation \ref{eq:alpha}) ) makes it possible to interpret disk fragmentation occurring at $t_{cool}\sim 
+\Omega^{-1}$ as the inability of the disk to sustain gravitational stress at $\alpha\gtrsim 1$ (Rice 2005)."
+ Iu this paper we investigate the structure aud evolution of exavitoturbuleut disks in which angular mionieutuui is transferred predominantly by the eravitational torques., In this paper we investigate the structure and evolution of gravitoturbulent disks in which angular momentum is transferred predominantly by the gravitational torques.
+" This problem has been previously investigated by Lin Pringle (1987) but with a rather naive prescription for the effective viscosity,", This problem has been previously investigated by Lin Pringle (1987) but with a rather naive prescription for the effective viscosity.
+ Also. some efforts have been devoted to understanding the structure of the eravitoturbulent disks which are unstable to frae1ieutation on large scale. ie. disks having Q~Qu and oe;~l evervwhere (Rafikoy 2005. 2007: Matznuer Levin 2005).," Also, some efforts have been devoted to understanding the structure of the gravitoturbulent disks which are unstable to fragmentation on large scale, i.e. disks having $Q\approx Q_0$ and $\alpha_{GI}\sim 1$ everywhere (Rafikov 2005, 2007; Matzner Levin 2005)."
+ Iu this work viscous evolution of the disk is explored according to the prescription (1)) without fiiug the value of acy instead it is calculated selt-consistently based on the plysical properties of the eas., In this work viscous evolution of the disk is explored according to the prescription \ref{eq:alpha}) ) without fixing the value of $\alpha_{GI}$ — instead it is calculated self-consistently based on the physical properties of the gas.
+ We concentrate our attention on rather cool disks in which opacity is due to dust erains thus focussing on the CIiu the outer parts of protostellar disks aud disks around supermassive black holes., We concentrate our attention on rather cool disks in which opacity is due to dust grains thus focussing on the GI in the outer parts of protostellar disks and disks around supermassive black holes.
+" We consider a eravitoturbulent disk iu which the dissipation of transient deusitv waves excited by GI is capable ofmaintaining Q=Qy. aud the cooling time f.,4: Is longer than οil"," We consider a gravitoturbulent disk in which the dissipation of transient density waves excited by GI is capable ofmaintaining $Q=Q_0$, and the cooling time $t_{cool}$ is longer than $\Omega^{-1}$."
+ Cooling time of the disk is Fin). where Xo ds the surface deusitv of the disk. e;=(kpTfp) is the isothermal sound. speed determined by the midplane temperature T. and f(7) is à functiou of the optical depth 7=fawpds (gp aud p are the eas opacity and deusitv. : is the vertical coordinate) which links the euütted fux F to T: F=oTlífir).," Cooling time of the disk is ), where $\Sigma$ is the surface density of the disk, $c_s\equiv (k_B T/\mu)$ is the isothermal sound speed determined by the midplane temperature $T$ , and $f(\tau)$ is a function of the optical depth $\tau=\int \kappa\rho dz$ $\kappa$ and $\rho$ are the gas opacity and density, $z$ is the vertical coordinate) which links the emitted flux $F$ to $T$: $F=\sigma T^4/f(\tau)$."
+ A specific form of f(r) depends on the wav in which energv is transported from the midplane of an optically thick disk to its plotosphere where it is radiated to space., A specific form of $f(\tau)$ depends on the way in which energy is transported from the midplane of an optically thick disk to its photosphere where it is radiated to space.
+ Rafikov (2007) has calculated f(r) in the case of efficicutly convecting disks., Rafikov (2007) has calculated $f(\tau)$ in the case of efficiently convecting disks.
+ Towever. in this work we assune (as was previously done in Rafikov 2005) for siuuplicity that euergv is carried from the disk micplane to its surface solely by radiation du which case f(rz) can be reasonably well approximated by This expression smoothlv interpolates between the cooling rates applicable in the optically thick (72 1) and optically thin (7« 1) regimes.," However, in this work we assume (as was previously done in Rafikov 2005) for simplicity that energy is carried from the disk midplane to its surface solely by radiation in which case $f(\tau)$ can be reasonably well approximated by This expression smoothly interpolates between the cooling rates applicable in the optically thick $\tau\gg 1$ ) and optically thin $\tau\ll 1$ ) regimes."
+" We asstunme a tempcrature-dependeut opacity in the form. kappa,τς which is appropriate at low temperatures when A ds dominated by dust grains."," We assume a temperature-dependent opacity in the form _0, which is appropriate at low temperatures when $\kappa$ is dominated by dust grains."
+ At very low temperatures. T«150 K. it is generally found (Bell Lin 1991: Semenov 2003) that opacity is duc to the icy eraius ancl is characterized by )j—22 3.IK. within a factor of 2 or so.," At very low temperatures, $T<150$ K, it is generally found (Bell Lin 1994; Semenov 2003) that opacity is due to the icy grains and is characterized by 2 ^2 within a factor of 2 or so."
+ At higher temperatures ices evaporate aud opacity behavior cau be crudely deseribed ase cmOLLT7? one e1 (Bell Lin 1991)., At higher temperatures ices evaporate and opacity behavior can be crudely described as $\kappa\approx 0.1T^{1/2}$ $^2$ $^{-1}$ (Bell Lin 1994).
+ For simplicity in this work we do not distinguish between the Rossclanc mean aud the Planck mean opacities (appropriate for TOlaud rcl correspondingly) as they lave similar values at low T., For simplicity in this work we do not distinguish between the Rosseland mean and the Planck mean opacities (appropriate for $\tau\gg 1$ and $\tau\ll 1$ correspondingly) as they have similar values at low $T$.
+ Now. using definition (2)) aud condition Q=Quy we fiud that in a gravitoturbuleut disk.," Now, using definition \ref{eq:Q}) ) and condition $Q=Q_0$ we find that )^2 in a gravitoturbulent disk."
+ Iu the optically thick regine total optical depth is dominated bv the midplane lavers of the disk i which most of the mass is conceutrated. so that up to factors of order unitv rz&(T)X.," In the optically thick regime total optical depth is dominated by the midplane layers of the disk in which most of the mass is concentrated, so that up to factors of order unity $\tau\approx \kappa(T)\Sigma$ ."
+ Clearly. this approxinatiou also works iu the optically thin case.," Clearly, this approximation also works in the optically thin case."
+ Thus. using equation (12)) one rather generally fiuds that 2," Thus, using equation \ref{eq:T}) ) one rather generally finds that ."
+"1, We can also calculate o¢y characterizing aneularmonentuni transport caused by the nou-axisviuuetric surface density perturbations.", We can also calculate $\alpha_{GI}$ characterizing angularmomentum transport caused by the non-axisymmetric surface density perturbations.
+ Using equations (1)). (6)). and (123) one fluds that Theὑ kinematicM viscosityM v=acerJο) isBthen given+ by the following expression:," Using equations \ref{eq:alpha}) ), \ref{eq:t_cool}) ), and \ref{eq:T}) ) one finds that The kinematic viscosity $\nu\equiv \alpha_{GI}c_s^2/\Omega$ isthen given by the following expression:"
+ ITustituto de Astronomiaa. Universidad Nacional Autóuuonia de Méxxico. Apdo.,"$^{1}$ Instituto de a, Universidad Nacional Autónnoma de Méxxico, Apdo."
+ Postal 70261L. Cd.," Postal 70–264, Cd."
+ Universitaria. Méxxico D.F. 01510 ? Copernicus Astronomical Contre. ul," Universitaria, Méxxico D.F. 04510 $^{2}$ Copernicus Astronomical Centre, ul."
+" Dartycka 18. 00.716 Warszawa. Poland ""Dustitute of Astronomy. Zielona (ιόντα University. ul."," Bartycka 18, 00–716 Warszawa, Poland $^{3}$ Institute of Astronomy, Zielona Górra University, ul."
+" Lubuska 2. 65265 Ziclona (ιόντα, ‘Department of Physics. University of Chicago. 5610 S. Ellis Ave.."," Lubuska 2, 65–265 Zielona Górra, $^{4}$ Department of Physics, University of Chicago, 5640 S. Ellis Ave.,"
+ Chicago. IL 60637 eauail: wvleecastroscuunuinnu.mx. wlodekocanmk.edu.pl. jduix(«uchicago.edu In this paper we studs the binary coalescence of a black hole and a quark star.," Chicago, IL 60637 e-mail: wleeastroscu.unam.mx, wlodekcamk.edu.pl, jdnixuchicago.edu In this paper we study the binary coalescence of a black hole and a quark star."
+ Stellar population studies indicate that if quark stars and. black holes exist at all. such binaries should exist in numbers significant from. the point of view of next-generation laser interferometric gravitational wave detectors. but smaller than the number of Lulse-Favlor type binaries (Belezvisskict al.," Stellar population studies indicate that if quark stars and black holes exist at all, such binaries should exist in numbers significant from the point of view of next-generation laser interferometric gravitational wave detectors, but smaller than the number of Hulse-Taylor type binaries (Belczyńsski et al."
+ 2001)., 2001).
+ C'oalescing quark stars also remain strong candidates, Coalescing quark stars also remain strong candidates
+of the formation of Alilkyw Way-sizecl halos to shed. more light on these questions. in. particular. by investigating a varicty of feedback. processes known to be important. in galaxy formation.,"of the formation of Milky Way-sized halos to shed more light on these questions, in particular by investigating a variety of feedback processes known to be important in galaxy formation."
+ Desides the impact of reionization. these include galaetic winds and outllows. energy input by growing supermassive black holes. or the non-thermal support ofgas by cosmic ravs or magnetic Gelds.," Besides the impact of reionization, these include galactic winds and outflows, energy input by growing supermassive black holes, or the non-thermal support of gas by cosmic rays or magnetic fields."
+ Ultimately we aim to reach similar numerical resolution as has been obtained for recent collisionless simulations. even though this goal may still be several vears. away.," Ultimately we aim to reach similar numerical resolution as has been obtained for recent collisionless simulations, even though this goal may still be several years away."
+ In this work. we present some of our first results.," In this work, we present some of our first results."
+ We use several well resolved hyvdrocdynamical simulations of the formation of a Alilky Way sized galaxy to investigate the properties of the predicted population of satellite galaxies. for different choices of the included physics.," We use several well resolved hydrodynamical simulations of the formation of a Milky Way sized galaxy to investigate the properties of the predicted population of satellite galaxies, for different choices of the included physics."
+ Besides a default reference model that includes only a treatment of radiative cooling. star formation. and cosmic reionization. we consider also models that add galactic winds. supermassive black hole growth. or cosmic ray injection by supernovae shock waves.," Besides a default reference model that includes only a treatment of radiative cooling, star formation, and cosmic reionization, we consider also models that add galactic winds, supermassive black hole growth, or cosmic ray injection by supernovae shock waves."
+ By comparing the simulation results with a comprehensive catalogue of the known Alilkw Way satellites. we seek. to determine which of these processes. is most. important. in shaping the satellite population.," By comparing the simulation results with a comprehensive catalogue of the known Milky Way satellites, we seek to determine which of these processes is most important in shaping the satellite population."
+ This paper is organized as follows., This paper is organized as follows.
+ In Section. ??.. we cleseribe the methodological details of our simulations. while the observational knowledge about the satellites is briefly summarized in Section ??..," In Section \ref{sec:methodology}, we describe the methodological details of our simulations, while the observational knowledge about the satellites is briefly summarized in Section \ref{sec:observations}. ."
+ Sections οον 2? and 77. present the results for our simulated populations of satellite galaxies. both with respect to individual satellite histories as well as with respect to their population as a whole.," Sections \ref{sec:abundance}, \ref{sec:history} and \ref{sec:scaling} present the results for our simulated populations of satellite galaxies, both with respect to individual satellite histories as well as with respect to their population as a whole."
+ Qur conclusions are summarized in Section ?7.., Our conclusions are summarized in Section \ref{sec:conclusions}.
+ Our. simulations are based on initial conditions originally constructed for the Aquarius Project (2). of the Virgo Consortium., Our simulations are based on initial conditions originally constructed for the Aquarius Project \citep{Aquarius} of the Virgo Consortium.
+ This project carried out highly resolved dark matter only simulations of 6 cifferent Alilkw Wav-sized valos. at a variety of cüfferent numerical resolutions.," This project carried out highly resolved dark matter only simulations of 6 different Milky Way-sized halos, at a variety of different numerical resolutions."
+" In the nomenclature of ?.. the halo and resolution level investigated vere ds called ""οι and has about 5.4 million. dark matter particles in the final virial radius."," In the nomenclature of \citet{Aquarius}, the halo and resolution level investigated here is called `Aq-C-4', and has about 5.4 million dark matter particles in the final virial radius."
+ Phe same object vas also been studied in the hydrodvnamie simulations of 7.. albeit at the considerably lower resolution (bv a factor S in particle number) corresponding to q-6C-5'.," The same object has also been studied in the hydrodynamic simulations of \citet{Scannapieco2009}, albeit at the considerably lower resolution (by a factor 8 in particle number) corresponding to `Aq-C-5'."
+ Phere it was ound that this -C’-halo produced the lowest bulge-to-disk ratio among the 6 candidate halos selected in the Aquarius >roject. making it a particularly σου candidate for the ormation of a Lage disk galaxy.," There it was found that this `C'-halo produced the lowest bulge-to-disk ratio among the 6 candidate halos selected in the Aquarius Project, making it a particularly good candidate for the formation of a large disk galaxy."
+ We note. however. that we do not expect our choice of target halo to influence our wincipal conclusions for the satellite population.," We note, however, that we do not expect our choice of target halo to influence our principal conclusions for the satellite population."
+ In this paper. we model the gas component with smoothed particle hyvdrodynamies (SPILL) and. introduce he gas particles into the initial conditions bv. splitting cach original particle into a dark matter and. gas. particle oir. displaced slightly. with respect. to cach other (at ixecl center-of-mass) to arrive at a regular distribution of the mean particle separations. and with a mass ratio corresponding to a baryon fraction of 16 per cent.," In this paper, we model the gas component with smoothed particle hydrodynamics (SPH) and introduce the gas particles into the initial conditions by splitting each original particle into a dark matter and gas particle pair, displaced slightly with respect to each other (at fixed center-of-mass) to arrive at a regular distribution of the mean particle separations, and with a mass ratio corresponding to a baryon fraction of 16 per cent."
+" The cosmological parameters. are ©,=0.25. O4=0.75. ax=0.9 and h=0.73. the same ones used. as in the original Aquarius simulations. which are consistent with the WALAPI cosmological constraints."," The cosmological parameters are $\Omega_m =
+0.25$, $\Omega_{\Lambda}=0.75$, $\sigma_8=0.9$ and $h=0.73$, the same ones used as in the original Aquarius simulations, which are consistent with the WMAP1 cosmological constraints."
+" A periodic box of size 1005!Mpe on a side is simulated. with varving spatial resolution that ""zooms in’ on the formation of a single galaxy."," A periodic box of size $100\,h^{-1}{\rm Mpc}$ on a side is simulated, with varying spatial resolution that `zooms in' on the formation of a single galaxy."
+ In the high-resolution region. we reach a mass simulation of =2105!M. and zm210hHAL. for dark matter and gas particles. respectively.," In the high-resolution region, we reach a mass simulation of $\approx 2\times
+10^5 \,h^{-1}{\rm M_{\odot}}$ and $\approx 2\times 10^4 \,h^{-1}{\rm
+ M_{\odot}}$ for dark matter and gas particles, respectively."
+ A constant comoving gravitational softening length of «=0.25.tkpe was use for all high-resolution particles.," A constant comoving gravitational softening length of $\epsilon = 0.25 \,
+h^{-1} \mathrm{kpc}$ was used for all high-resolution particles."
+ We employed the parallel PreeSPLE code for our runs. which is an improved. and. extended version of (?).," We employed the parallel TreeSPH code for our runs, which is an improved and extended version of \citep{Gadget2}."
+ caleulates the long-range gravitationa fick in. Fourier space. and the short range forces in rea space with a hierarchical multipole expansion. based on a tree.," calculates the long-range gravitational field in Fourier space, and the short range forces in real space with a hierarchical multipole expansion, based on a tree."
+ This approach guarantees a homogeneously high spatia resolution in the gravitational force caleulation and can be ellicientlv combined with an individual timestep integration scheme., This approach guarantees a homogeneously high spatial resolution in the gravitational force calculation and can be efficiently combined with an individual timestep integration scheme.
+ For the hydrodynamiecs. uses the “entropy formulation’ of SPLI (2). which is derived from a variational principle and simultaneously. conserves energy. and entropy where appropriate.," For the hydrodynamics, uses the `entropy formulation' of SPH \citep{Springel2002}, which is derived from a variational principle and simultaneously conserves energy and entropy where appropriate."
+ In the hydrocdynamic part ofADGIT.. different physical processes besides. ordinary gas dynamics are calculated.," In the hydrodynamic part of, different physical processes besides ordinary gas dynamics are calculated."
+ Alost importantly. these are radiative cooling. star formation and its regulation by supernovae feedback. processes. (?)..," Most importantly, these are radiative cooling, star formation and its regulation by supernovae feedback processes \citep{SFR_paper}."
+ The code can optionally also mocel black hole growth and (7) and cosmic rav physics (2)..., The code can optionally also model black hole growth and \citep{BH_paper} and cosmic ray physics \citep{CR_Paper}.
+ We shall now. brielly deseribe the physics modules we used., We shall now briefly describe the physics modules we used.
+ ltadiative cooling is followed. for a primordial. mixture of helium. anc hydrogen. under the assumption of collisional ionization equilibrium. using a formulation as in 7..," Radiative cooling is followed for a primordial mixture of helium and hydrogen under the assumption of collisional ionization equilibrium, using a formulation as in \citet{Katz1996}."
+ A spatially uniform. ionizing UV. background. is. introduced with an amplitude and. time evolution described by an updated version of ?.. leading to reionization of the universe w redshift 26.," A spatially uniform, ionizing UV background is introduced with an amplitude and time evolution described by an updated version of \citet{Haardt1996}, leading to reionization of the universe by redshift $z\simeq 6$."
+ To model star formation. we use the hybrid multiphase moclel for star formation and supernova feedback introduced w 7. dn which every sullicicntly dense gas particle is treated as a representative region for the multiphase structure of he interstellar medium (ISM).," To model star formation, we use the hybrid multiphase model for star formation and supernova feedback introduced by \citet{SFR_paper}, in which every sufficiently dense gas particle is treated as a representative region for the multiphase structure of the interstellar medium (ISM)."
+ These hybrid. particles are λογος to be comprised. of cold. dense elouds in. pressure equilibrium with a hot ambient gas. phase. where only he clouds contribute material available for star formation.," These hybrid particles are pictured to be comprised of cold dense clouds in pressure equilibrium with a hot ambient gas phase, where only the clouds contribute material available for star formation."
+ Alass and energy exchange processes between these two ohases are computed: by simple differential equations. as deseribed. in. 2.. giving rise to an effective. equation. of state that regulates the dense gas of the ISM.," Mass and energy exchange processes between these two phases are computed by simple differential equations, as described in \citet{SFR_paper}, giving rise to an effective equation of state that regulates the dense gas of the ISM."
+ Collisionless star particles are spawned stochastically from this. star-forming phase according to a local estimate of the star formation rate., Collisionless star particles are spawned stochastically from this star-forming phase according to a local estimate of the star formation rate.
+ Phe gas consumption timescale of the mocel is calibrated such that itreproduces the Kennicutt law (?) between star formation rate and gas surface density observed in low-redshift disk galaxies.," The gas consumption timescale of the model is calibrated such that itreproduces the Kennicutt law \citep{Kennicutt}
+ between star formation rate and gas surface density observed in low-redshift disk galaxies."
+work quite well. the eror being always less than We show now how it is possible to use the simple technique of the inverse problemi to estimate the slope a of the lens AIF and the dark halo uass fraction f composed bv MACIIOs for each one of the imocdels labelled bv the codes previously explained.,"work quite well, the error being always less than We show now how it is possible to use the simple technique of the inverse problem to estimate the slope $\alpha$ of the lens MF and the dark halo mass fraction $f$ composed by MACHOs for each one of the models labelled by the codes previously explained."
+ We consider data comine from the first 5.7 vears of observations towards LMC by the MACTIO collaboration (Alcocketal. 2000a)). limiting ourselves to the thirtecu eveuts selected according to the ccalledA.," We consider data coming from the first 5.7 years of observations towards LMC by the MACHO collaboration \cite{A00}) ), limiting ourselves to the thirteen events selected according to the called."
+ These are high S/N events and are spatially distributed i a way which is cousistcut with the hypothesis that they are due to lenses belongiug to our halo and not to LMC self lensine: we will discuss of this problem later ou., These are high S/N events and are spatially distributed in a way which is consistent with the hypothesis that they are due to lenses belonging to our halo and not to LMC self lensing; we will discuss of this problem later on.
+ For this set of events we lave where <trDp IS πρίν the average value of the duration of the observed events., For this set of events we have where $< t_E >_{obsd}$ is simply the average value of the duration of the observed events.
+" As a test of the correctuess of this estimate we iav note that (7/LE)trDoug (with E=6.12«10* vvenrs) is exactly equal το TopsNobel, as it has to be for the reasous we are eoing to explain later."," As a test of the correctness of this estimate we may note that $(\pi/4E) < t_E >_{obsd}$ (with $E = 6.12 \times 10^{7}$ years) is exactly equal to $\tau_{obsd}/N_{ev}^{obsd}$, as it has to be for the reasons we are going to explain later."
+ We have then to estimate the uncertainties ou the observed quautities., We have then to estimate the uncertainties on the observed quantities.
+ First. we consider the directly observed optical depth. aud we simply use a method as similar as possible to the oue proposed by the ΔΙΑΠΟ eroup itself for a very conservative estinate of the error ou Teas (Alcocketal. 1997a)): we divide the observed events according to their duration frg in biu of 10 davs: im such a wav Tops IS more or less the same for events in the same bin aud the errors are approximately polssoniau.," First, we consider the directly observed optical depth, and we simply use a method as similar as possible to the one proposed by the MACHO group itself for a very conservative estimate of the error on $\tau_{meas}$ \cite{A97a}) ): we divide the observed events according to their duration $t_E$ in bin of 10 days; in such a way $\tau_{obsd}$ is more or less the same for events in the same bin and the errors are approximately poissonian."
+" For cach bin we estimate Nyon=Neesbin)VNebsd(bin) and N,,=Nbhin)lyNebsd(bin) and define τι and r(""""* as the iain and naxiuuu value of τι for the events iu that bin (being 7=ate/ LE).", For each bin we estimate $N_{low} = N_{ev}^{obsd}(bin) - \sqrt{N_{ev}^{obsd}(bin)}$ and $N_{up} = N_{ev}^{obsd}(bin) + \sqrt{N_{ev}^{obsd}(bin)}$ and define $\tau_{1}^{min}$ and $\tau_{1}^{max}$ as the minimum and maximum value of $\tau_{1}$ for the events in that bin (being $\tau_1 = \pi t_E / 4 E$ ).
+ Then we estimate At the cud we fiud The error ou the optical depth turus out to be so lavee (~ YO) because of the limited πιο of eveuts., Then we estimate At the end we find The error on the optical depth turns out to be so large $\sim 70 \%$ ) because of the limited number of events.
+ Let us turn now to the error on the nunber of observed events., Let us turn now to the error on the number of observed events.
+ This simply comes from the low statistics and may be assumed to be poissonian. ie. AVEVers’líV A.," This simply comes from the low statistics and may be assumed to be poissonian, i.e. $\delta N_{ev}^{obsd}/ N_{ev}^{obsd} = 1/\sqrt{N_{ev}^{obsd}} \simeq 28\%$ ."
+ Finally. the error ou «fgE> is obtained by propagating the error ou Nobel which gives us The observed quantities 1nav be now compared to the theoretical oues evaluated in the previous section. which also take iuto account the detection efücienev du order to make the comparison meaninetul.," Finally, the error on $< t_E >$ is obtained by propagating the error on $N_{ev}^{obsd}$ which gives us The observed quantities may be now compared to the theoretical ones evaluated in the previous section, which also take into account the detection efficiency in order to make the comparison meaningful."
+ To do this we mist first remember that the theoretical quautitics have been calculated under the hwpothesis that the dark halo is totally made out by ATACTIOs. ie. with f=1.," To do this we must first remember that the theoretical quantities have been calculated under the hypothesis that the dark halo is totally made out by MACHOs, i.e. with $f = 1$."
+ Actually. the exact value of f is not well known: from the more recent observational constraints it is quite unlikely that f=l.," Actually, the exact value of $f$ is not well known: from the more recent observational constraints it is quite unlikely that $f = 1$."
+ However to take iuto account of f is not very ifhcult: we have simply to multiply by f£ the expression of the differential rate dUfete aud consequently the ones obtained for NUM aul Tae., However to take into account of $f$ is not very difficult: we have simply to multiply by $f$ the expression of the differential rate $d\Gamma/dt_E$ and consequently the ones obtained for $N_{ev}^{oble}$ and $\tau_{oble}$ .
+ Note that <fp> ds Ecrdepeudent ou the value of £., Note that $< t_E >$ is independent on the value of $f$.
+" Then we have the following clatious between observable aud observed quantities: Dividing these two equatious and using the relation TomeiNOM=(nÍLIE)«te one Bets: Taps,V=(TALE)<fg μμ. Which nav be used to test the correctuess of our previous estimate of «xteDb."," Then we have the following relations between observable and observed quantities: Dividing these two equations and using the relation $\tau_{oble}/N_{ev}^{oble} = (\pi/4E)< t_E >$, one gets: $\tau_{obsd}/N_{ev}^{obsd} = (\pi/4E)< t_E >_{obsd}$ , which may be used to test the correctness of our previous estimate of $< t_E >_{obsd}$."
+" For each model. we nav solve the system (27)) im the ""uknowns o and f aud estimate them toghether with the errors connected to our analysis."," For each model, we may solve the system \ref{eq: system}) ) in the unknowns $\alpha$ and $f$ and estimate them toghether with the errors connected to our analysis."
+" Solving Eqs.(27)) is not really useful since the very ligh error ou 7,44; leads us to ect 10 constraints at all on the parameters (a.f£)."," Solving \ref{eq: system}) ) is not really useful since the very high error on $\tau_{obsd}$ leads us to get no constraints at all on the parameters $(\alpha, f)$."
+ We may then use a third relation we have at our disposal. given bx which is quite easv to solve nunuericallv to get au estimate of the slope à of the ME.," We may then use a third relation we have at our disposal, given by which is quite easy to solve numerically to get an estimate of the slope $\alpha$ of the MF."
+ Solving Eq.(28)). taking iuto account also the uncertainties. will eive us in general more than one solution compatible with the mucrolensing data.," Solving \ref{eq: tecompare}) ), taking into account also the uncertainties, will give us in general more than one solution compatible with the microlensing data."
+ A further selection can be done imposing that a must be in the rauge (0.0.5.0) since values outside this range ave reasonably quiteunlikely.," A further selection can be done imposing that $\alpha$ must be in the range $(0.0, 5.0)$ since values outside this range are reasonably quite."
+".. Iu this wav. for cach model. we have estimated a rauge (αμαν} for the slope of the ME simply requiring that being ófg/fr the fractional error on <tr5 Which we have previously estimated to be of order of Iaving estimated à. and beingmm AWeeTee known, we may how ect a constraint also on f: ou the plot PNght as a function of a and f we make the contour levels for NO_NotedON=0.39 and EN_NU|283%= 16.61."," In this way, for each model, we have estimated a range $(\alpha_l, \alpha_u)$ for the slope of the MF simply requiring that being $\delta t_E/ t_E$ the fractional error on $< t_E >_{obsd}$ which we have previously estimated to be of order of Having estimated $\alpha$ and being $N_{ev}^{obsd}$ known, we may now get a constraint also on $f$: on the plot $f N_{ev}^{oble}$ as a function of $\alpha$ and $f$ we make the contour levels for $f N_{ev}^{oble} = N_{ev}^{obsd} - 28\,\% = 9.39$ and $f N_{ev}^{oble} = N_{ev}^{obsd} + 28\,\% = 16.61$ ."
+" On this eraph one has to add also the vertical lines corresponding to a=a; aud a=a, ({see Fig.", On this graph one has to add also the vertical lines corresponding to $\alpha = \alpha_l$ and $\alpha = \alpha_u$ (see Fig.
+ 2)., 2).
+ The region R of the parameter space (a.f) delimited by the two level curves aud these two vertical ines is that which is consistent with the constraints on a and the ones coming from nücroleusiug observations owards LAIC. i.e. the nmuiber of observed events aud their uean duration.," The region $\cal{R}$ of the parameter space $(\alpha, f)$ delimited by the two level curves and these two vertical lines is that which is consistent with the constraints on $\alpha$ and the ones coming from microlensing observations towards LMC, i.e. the number of observed events and their mean duration."
+ From Fig., From Fig.
+ 2 one sees that itis possible to, 2 one sees that itis possible to
+The inunediate euvironiment of an accreting supermassive black hole is extremely exotic.,The immediate environment of an accreting supermassive black hole is extremely exotic.
+ Droad iron lines provide us with the best tool to date for studying these regious., Broad iron lines provide us with the best tool to date for studying these regions.
+ audX observatious have already shown us that the accretion disk iu at least some AGN extends very close to the black hole (aud mavbe so close as to sugeest that the black hole umst be rotating)., and observations have already shown us that the accretion disk in at least some AGN extends very close to the black hole (and maybe so close as to suggest that the black hole must be rotating).
+ Furthermore. the detection of broad iron line variability by is imos likely tracking structural changes in the accretion disk and/or N-ray euütting corona.," Furthermore, the detection of broad iron line variability by is most likely tracking structural changes in the accretion disk and/or X-ray emitting corona."
+ However. large effective area detectws are required to male further progress.," However, large effective area detectors are required to make further progress."
+WALA will alow these structural chauges to be characterized in detail. thereby xobiug the instabilities that affect the immer accretion disk/coroua.," will allow these structural changes to be characterized in detail, thereby probing the instabilities that affect the inner accretion disk/corona."
+" Furthermore.XAZAZ will allow us to study iron line variability caused by the accretion clisk rotatioi. allowing us Oo nuüeasure the ass of the dack hole aud constrain the location,litenue of the N-rayv flares."," Furthermore, will allow us to study iron line variability caused by the accretion disk rotation, allowing us to measure the mass of the black hole and constrain the location/lifetime of the X-ray flares."
+ Eventually.Hation-X will allow us to search for iron ine reverberation.," Eventually, will allow us to search for iron line reverberation."
+ The detection of reverevation will eive robust signatures of blacx hole spin iux provide the tools to study he inner disk structure m unprecedoenutec cletail., The detection of reverberation will give robust signatures of black hole spin and provide the tools to study the inner disk structure in unprecedented detail.
+ Further in the future. direct imaging of the iuer disk aid black hole region ln neambv ACN will be possible using N-rav inerferometry.," Further in the future, direct imaging of the inner disk and black hole region in nearby AGN will be possible using X-ray interferometry."
+ This will provide the ultimae observational probe of black hole astrophysics., This will provide the ultimate observational probe of black hole astrophysics.
+Ryutovaetal.(2008) proposed that some of the chromospheric brightenings are produced by shocks resulted from a sling-shot effect associated with magnetic reconnection process in neighboring penumbral filaments.,\cite{ryutova2008} proposed that some of the chromospheric brightenings are produced by shocks resulted from a sling-shot effect associated with magnetic reconnection process in neighboring penumbral filaments.
+ The velocity stratification shown in Fig., The velocity stratification shown in Fig.
+ 6 (b) have a possibility to provide useful information to know where magnetic reconnection takes place if the downflow is driven by this mechanism., \ref{vprof} (b) have a possibility to provide useful information to know where magnetic reconnection takes place if the downflow is driven by this mechanism.
+ The downflow in the lower photosphere suggests that a reconnection site might be located in the photosphere because significant downflows ought to be observed even in the upper and middle photosphere if magnetic reconnection occurs in the chromosphere., The downflow in the lower photosphere suggests that a reconnection site might be located in the photosphere because significant downflows ought to be observed even in the upper and middle photosphere if magnetic reconnection occurs in the chromosphere.
+" Because we used only three nodes along the optical depth in the Stokes inversion and derived the velocity stratification among the nodes with the interpolation, we cannot argue acceleration or deceleration of the downward flow from the inversion."," Because we used only three nodes along the optical depth in the Stokes inversion and derived the velocity stratification among the nodes with the interpolation, we cannot argue acceleration or deceleration of the downward flow from the inversion."
+ Further spectroscopic studies are necessary to resolve the height dependence of their flow structures from the middle to the upper photosphere., Further spectroscopic studies are necessary to resolve the height dependence of their flow structures from the middle to the upper photosphere.
+" Another possibility of the patchy downward flows are indications of convective rolls in penumbral filaments as proposed by Zakharovetal.(2008) who found blueshifts on the limbward side of a penumbral filament, and weak red shifts on the centerward side."," Another possibility of the patchy downward flows are indications of convective rolls in penumbral filaments as proposed by \cite{zakharov2008} who found blueshifts on the limbward side of a penumbral filament, and weak red shifts on the centerward side."
+" It is, however, not clear how the convective rolls make sporadic downflow patches."," It is, however, not clear how the convective rolls make sporadic downflow patches."
+ Ortizetal.(2010) found small concentrations of downflows associated with strong upflows in umbral dots., \cite{ortiz2010} found small concentrations of downflows associated with strong upflows in umbral dots.
+" Their sizes smaller than 0.5"" and lifetimes shorter than a few minutes look similar with the nature of the downflow patches in the penumbra studied in this paper.", Their sizes smaller than 0.5” and lifetimes shorter than a few minutes look similar with the nature of the downflow patches in the penumbra studied in this paper.
+" There is a possibility that magnetoconvection in the presence of strong magnetic fields transiently makes small concentration of downflows in the lower photosphere, which can be explored with numerical simulations recently developed (e.g.Rempeletal., 2009).."," There is a possibility that magnetoconvection in the presence of strong magnetic fields transiently makes small concentration of downflows in the lower photosphere, which can be explored with numerical simulations recently developed \citep[e.g.][]{rempel2009}. ."
+"When the Edgeworth series is truncated at linear order in op, the resulting pdf is given by p(v)) = where v=δµ/σῃ as before.","When the Edgeworth series is truncated at linear order in $\sigma_R$, the resulting pdf is given by ) = where $\nu\equiv \delta_R/\sigma_R$ as before."
+" Observe that if S3>0, a sufficiently large negative v gives p(v)<0 and, similarly, if S3< 0, a sufficiently large positive v gives p(v)«0."," Observe that if $ S_3>0$, a sufficiently large negative $\nu$ gives $p(\nu)<0$ and, similarly, if $S_3<0$ , a sufficiently large positive $\nu$ gives $p(\nu)<0$."
+" For instance, if |S3|=0.1, p(v) becomes negative as early as ~3."," For instance, if $|S_3|=0.1$, $p(\nu)$ becomes negative as early as $|\nu|\simeq3$."
+" This implies that a linear truncation of the Edgeworth series is highly suspect and is certainly not suitable for calculating,|v| for instance, the mass function whereby high values of density fluctuations are involved."," This implies that a linear truncation of the Edgeworth series is highly suspect and is certainly not suitable for calculating, for instance, the mass function whereby high values of density fluctuations are involved."
+ The Edgeworth series truncated at quadratic order in or yields p»)- We wish to determine the combination of $3 and $4 such that p(v) is non-negative., The Edgeworth series truncated at quadratic order in $\sigma_R$ yields ) = We wish to determine the combination of $ S_3$ and $ S_4$ such that $p(\nu)$ is non-negative.
+ The numerical evaluation of p(v) over a grid of S5 and S4 is shown in figure B]., The numerical evaluation of $p(\nu)$ over a grid of $S_3$ and $ S_4$ is shown in figure \ref{scatter}.
+" In the figure, we mark points for which p(v)>0 in the (on$3,07,94) plane, in the domain v€[--20,20]."," In the figure, we mark points for which $p(\nu)>0$ in the $(\sigma_R S_3,\sigma_R^2 S_4)$ plane, in the domain $\nu\in[-20,20]$."
+ This reveals a closed region in which p(v)>0., This reveals a closed region in which $p(\nu)>0$.
+" In fact, the bounding envelope can be found analytically by settingp(v)=p'(v)0, but the resulting equation has a very complicated parametric form which we shall not show here."," In fact, the bounding envelope can be found analytically by setting$p(\nu)=p^\pr(\nu)=0$, but the resulting equation has a very complicated parametric form which we shall not show here."
+ For details of this technique see ?.., For details of this technique see \cite{jondeau}.
+" A curious feature of figure is that the excess kurtosis, 7.54, is limited to a small, non-negative range."," A curious feature of figure \ref{scatter} is that the excess kurtosis, $\sigma^2 S_4$, is limited to a small, non-negative range."
+ We can prove this as follows., We can prove this as follows.
+" Setting [3]$3—0, the quadratic series can be written as ==14 This expression clearly achieves the minimum when 7?—3=0."," Setting $S_3=0$, the quadratic series can be written as = This expression clearly achieves the minimum when $\nu^2-3=0$."
+" Requiring the minimum to be non-negative establishes the upper bound 07S,<4.", Requiring the minimum to be non-negative establishes the upper bound $\sigma^2S_4\leq4$.
+" Next, if S4<0, the quartic expression is unbounded from below and so the pdf will be negative for some large x."," Next, if $S_4<0$, the quartic expression is unbounded from below and so the pdf will be negative for some large $x$."
+" Thus, we must have 0€0754<4 The bound for oS3 is more difficult to establish and we shall not go into the detail here."," Thus, we must have $0\leq\sigma^2 S_4\leq4$ The bound for $\sigma S_3$ is more difficult to establish and we shall not go into the detail here."
+" We simply note since an analytic expression describing the shaded region in figure | exists, the region in fact represents combinations of $3 and S4 νι, and gnz) for which the pdf is non-negative on the entire real line and not just in [-20,20]."," We simply note since an analytic expression describing the shaded region in figure \ref{scatter} exists, the region in fact represents combinations of $S_3$ and $S_4$ $\fnl$ and $\gnl$ ) for which the pdf is non-negative on the entire real line and not just in [-20,20]."
+" For higher-order truncations, it becomes increasingly difficult to find such a region, as expected given the conclusion in section ?? Figure shows the same set of axes as figure with the Edgeworth series now expanded up to terms of order c?, σὃ, o1? and[4) o?° (top row to bottom row)."," For higher-order truncations, it becomes increasingly difficult to find such a region, as expected given the conclusion in section \ref{sectionrose}
+ Figure \ref{scatterlots} shows the same set of axes as figure \ref{scatter} with the Edgeworth series now expanded up to terms of order $\sigma^3$, $\sigma^5$, $\sigma^{10}$ and $\sigma^{20}$ (top row to bottom row)."
+" In producing[3] these figures, we have set the rest of the cumulants to zero."," In producing these figures, we have set the rest of the cumulants to zero."
+" This is roughly equivalent to parametrizing the non-Gaussianity by fwr, and gwr only.", This is roughly equivalent to parametrizing the non-Gaussianity by $\fnl$ and $\gnl$ only.
+" Note that if n is odd, the series up to n terms performs significantly worse than one with even n."," Note that if $n$ is odd, the series up to $n$ terms performs significantly worse than one with even $n$ ."
+" This is simply because odd (Hermite) polynomials are not positive definite, whereas even ones are, provided the coefficients are properly chosen."," This is simply because odd (Hermite) polynomials are not positive definite, whereas even ones are, provided the coefficients are properly chosen."
+" When scanning over à sufficiently large range of v, an odd-ordered Edgeworth expansion will not produce any well-defined pdf whatsoever."," When scanning over a sufficiently large range of $\nu$, an odd-ordered Edgeworth expansion will not produce any well-defined pdf whatsoever."
+ 'The sensitivity of the regions to the range of v considered is clearly seen inthe difference between the column on the left (in which p(v) is only required to be non-negative for |v|« 5) and on the right (|v|< 20)., The sensitivity of the regions to the range of $\nu$ considered is clearly seen inthe difference between the column on the left (in which $p(\nu)$ is only required to be non-negative for $|\nu|<5$ ) and on the right $|\nu|<20$ ).
+" As the range of v increases, the cluster of points shrinks asit becomes increasingly difficult to find a closed region with p(v)> 0."," As the range of $\nu$ increases, the cluster of points shrinks asit becomes increasingly difficult to find a closed region with $p(\nu)>0$ ."
+"detail the degeneracies in the a—M"" plane.",detail the degeneracies in the $\alpha - M^*$ plane.
+ The results are shown in Figs., The results are shown in Figs.
+ 6 and 7.., \ref{fig:contours} and \ref{fig:mfix}.
+" In the first figure, we analysed the redshift intervals where both parameters were allowed to vary, while in the second one we studied the dependence of the α on the chosen M* in those redshift bins where we were forced to fix the characteristic mass to constrain the analysis."," In the first figure, we analysed the redshift intervals where both parameters were allowed to vary, while in the second one we studied the dependence of the best-fit $\alpha$ on the chosen $M^*$ in those redshift bins where we were forced to fix the characteristic mass to constrain the maximum-likelihood analysis."
+"In Fig. 6,,","In Fig. \ref{fig:contours},"
+" we show the lo and 2c contours for a and M* Schechter parameters, for both the BC03-based GSMFs (black solid curves) and CB07-based ones (red dotted curves)."," we show the $1\sigma$ and $2\sigma$ contours for $\alpha$ and $M^*$ Schechter parameters, for both the BC03-based GSMFs (black solid curves) and CB07-based ones (red dotted curves)."
+" While the parameter α is well-constrained at all redshifts (although with uncertainties increasing with z), our data prevent us from properly inferring the value of the characteristic mass."," While the parameter $\alpha$ is well-constrained at all redshifts (although with uncertainties increasing with $z$ ), our data prevent us from properly inferring the value of the characteristic mass."
+" Nonetheless, as we show below, the result on a is robust against the degeneracy of M*."," Nonetheless, as we show below, the result on $\alpha$ is robust against the degeneracy of $M^*$."
+ The steepening in a between z~0.8 and z~3 is clear from Fig., The steepening in $\alpha$ between $z\sim 0.8$ and $z\sim 3$ is clear from Fig.
+ 6 when the BCO3 stellar library is used., \ref{fig:contours} when the BC03 stellar library is used.
+" When we instead adopted CBO07 templates, the faint-end slope did not change much from z~0.8 to z~ 2.2, while at higher redshifts z« 3.5) we were forced to fix the value of M* to constrain the fit (see Sect. 3.4)),"," When we instead adopted CB07 templates, the faint-end slope did not change much from $z\sim 0.8$ to $z\sim 2.2$ , while at higher redshifts $2.5<z<3.5$ ) we were forced to fix the value of $M^*$ to constrain the fit (see Sect. \ref{sec:templates}) ),"
+ making the best-fit α parameter dependent on the choice of the characteristic mass., making the best-fit $\alpha$ parameter dependent on the choice of the characteristic mass.
+" To study how the best-fit value of a changes when varying the M* value, we built a grid of log(M"" [Ms]) ranging from 10.5 to 11.6 (these limits are justified by previous works) with steps in mass of 0.1 and we fitted a Schechter function to the data for each element of the grid."," To study how the best-fit value of $\alpha$ changes when varying the $M^*$ value, we built a grid of $M^*[M_\odot]$ ) ranging from 10.5 to 11.6 (these limits are justified by previous works) with steps in mass of 0.1 and we fitted a Schechter function to the data for each element of the grid."
+ We adopted this procedure in all the redshifts bins where the fit is unconstrained., We adopted this procedure in all the redshifts bins where the fit is unconstrained.
+ We show in Fig., We show in Fig.
+" 7 the a—M* plane, where the shaded areas show the values of the faint-end slope at different fixed M* in the redshift bins indicated by the labels."," \ref{fig:mfix} the $\alpha - M^*$ plane, where the shaded areas show the values of the faint-end slope at different fixed $M^*$ in the redshift bins indicated by the labels."
+" The symbols show the best-fit values (and their uncertainties) for « and M* at z«3.5 (z«2.5 when using CBO7 templates), where our results are insignificantly affected by the lack of high quality statisticsand we could allow both parameters to vary."," The symbols show the best-fit values (and their uncertainties) for $\alpha$ and $M^*$ at $z<3.5$ $z<2.5$ when using CB07 templates), where our results are insignificantly affected by the lack of high quality statisticsand we could allow both parameters to vary."
+" The upper panel refers to the BCO3-based GSMFs, while the lower one is obtained by adopting CBO7 stellar templates."," The upper panel refers to the BC03-based GSMFs, while the lower one is obtained by adopting CB07 stellar templates."
+" From Fig.7 (lower panel, blue shaded region), it appears that, whatever reasonable value for M* is chosen at 2.5«z« 3.5, the best-fit @ is clearly steeper than the best-fit values at lower redshifts, confirming the result found with BCO3 templates and supporting that the major result of this paper is unaffected by the poor sampling of the high mass regime."," From \ref{fig:mfix} (lower panel, blue shaded region), it appears that, whatever reasonable value for $M^*$ is chosen at $2.5<z<3.5$ , the best-fit $\alpha$ is clearly steeper than the best-fit values at lower redshifts, confirming the result found with BC03 templates and supporting that the major result of this paper is unaffected by the poor sampling of the high mass regime."
+" However, as shown by the red shaded regions in Fig. 7,,"," However, as shown by the red shaded regions in Fig. \ref{fig:mfix}, ,"
+ presently available data do not allow us to draw any firm conclusionabout the trend of o between z~3 and z~ 4., presently available data do not allow us to draw any firm conclusionabout the trend of $\alpha$ between $z\sim 3$ and $z\sim 4$ .
+" To corroborate our result for the slope of the low-mass end, and ensure that is unaffected by the lack of statistics at"," To corroborate our result for the slope of the low-mass end, and ensure that is unaffected by the lack of statistics at"
+In this section I will analyze the behavior of the main emission lines in the optical domain.,In this section I will analyze the behavior of the main emission lines in the optical domain.
+" For convenience, I will express the line fluxes inRayleighs?."," For convenience, I will express the line fluxes in."
+. The [OI]5577 is generally the most prominent feature in the optical night sky spectrum., The [OI]5577 is generally the most prominent feature in the optical night sky spectrum.
+" It falls right in the center of the V passband, giving a typical contribution of to the global surface brightness in this filter."," It falls right in the center of the $V$ passband, giving a typical contribution of to the global surface brightness in this filter."
+" It has a typical intensity of ~250 R, it arises in layers placed at about 90 km (Roach Gordon 1973)) and it displays a marked dependency on solar activity (Rayleigh 1928))."," It has a typical intensity of $\sim$ 250 R, it arises in layers placed at about 90 km (Roach Gordon \cite{roach}) ) and it displays a marked dependency on solar activity (Rayleigh \cite{rayleigh}) )."
+" This is clearly shown also by the data presented here (Fig. 9,,"," This is clearly shown also by the data presented here (Fig. \ref{fig:oi5577},"
+" lower panel), which indicate also the presence of pronounced fluctuations (40 to 750 R peak-to-peak) around the average level (~230 R)."," lower panel), which indicate also the presence of pronounced fluctuations (40 to 750 R peak-to-peak) around the average level $\sim$ 230 R)."
+" To quantify the correlation with solar activity and following the procedure that has been applied to the broad band data (see Sec. 4)),"," To quantify the correlation with solar activity and following the procedure that has been applied to the broad band data (see Sec. \ref{sec:sun}) ),"
+ I have fitted to the data a law of the type logF=Fo+ySFD.," I have fitted to the data a law of the type $\log F = \log F_0 + \gamma \;
+SFD$."
+" The results for this and other lines are presented in Table 6,, which shows also the number of data points used (NV) and the linear correlation coefficient (r)."," The results for this and other lines are presented in Table \ref{tab:lines}, which shows also the number of data points used $N$ ) and the linear correlation coefficient $r$ )."
+" What is interesting to note is that, once the solar dependency is removed from the data, the flux of this line displays a marked SAO (see Fig. 9,,"," What is interesting to note is that, once the solar dependency is removed from the data, the flux of this line displays a marked SAO (see Fig. \ref{fig:oi5577},"
+" upper panel), similar to that seen in the broad band data (Fig. 8))."," upper panel), similar to that seen in the broad band data (Fig. \ref{fig:season}) )."
+" This fact has been already noticed by Buriti et al. (2001)),"," This fact has been already noticed by Buriti et al. \cite{buriti}) ),"
+ who found that this line and other mesospheric features all show a SAO., who found that this line and other mesospheric features all show a SAO.
+" The same behavior is, in fact, shown by the doublet (see Fig. 10)),"," The same behavior is, in fact, shown by the doublet (see Fig. \ref{fig:oi6300}) ),"
+" which is produced at 250-300 km (Roach Gordon 1973)) and it is known to undergo abrupt intensity changes on two active regions about 20? on either side of the geomagnetic equator (Barbier 1957)), hence marginally including the Paranal site."," which is produced at 250-300 km (Roach Gordon \cite{roach}) ) and it is known to undergo abrupt intensity changes on two active regions about $^\circ$ on either side of the geomagnetic equator (Barbier \cite{barbier2}) ), hence marginally including the Paranal site."
+" Indeed, the [OI]6300 measured fluxes show very strong variations (10 to 950 R) around the average level (~150 R), with spikes reaching 1 ΚΚ. so that the line fluxes span almost a factor 100 (to be compared with the factor ~25 measured for the [OI]5577."," Indeed, the [OI]6300 measured fluxes show very strong variations (10 to 950 R) around the average level $\sim$ 150 R), with spikes reaching $\sim$ 1 kR, so that the line fluxes span almost a factor 100 (to be compared with the factor $\sim$ 25 measured for the [OI]5577."
+" In this respect, the behavior of the doublet is different from that of the [OI]5577 line, since in a significant number of cases its flux is very small and the line is practically invisible, lost in the OH(9-3) molecular band (see Fig. 1))."," In this respect, the behavior of the doublet is different from that of the [OI]5577 line, since in a significant number of cases its flux is very small and the line is practically invisible, lost in the OH(9-3) molecular band (see Fig. \ref{fig:skyspectrum}) )."
+" In fact, its flux is less than 80 R for more than of the cases considered in this work, the minimum recorded value being ~10 R (Fig. 11))."," In fact, its flux is less than 80 R for more than of the cases considered in this work, the minimum recorded value being $\sim$ 10 R (Fig. \ref{fig:linehist}) )."
+ The variation range of these two features and their relative contributions to the V and R total flux are summarized in Table 5.., The variation range of these two features and their relative contributions to the $V$ and $R$ total flux are summarized in Table \ref{tab:oi}.
+" From the values reported in this table, one can see that [OIJ5577 can produce a maximum variation of ~ 0.5 mag arcsec? in the V band, while a similar effect is produced by the [OI] doublet in the Rpassband‘."," From the values reported in this table, one can see that [OI]5577 can produce a maximum variation of $\sim$ 0.5 mag $^{-2}$ in the $V$ band, while a similar effect is produced by the [OI] doublet in the $R$."
+" Given the fact that the RMS variation in the V and R passbands is about 0.25 mag arcsec? (see Table 3)), this implies that the fluctuations seen in these passband are not completely accounted by the changes in the atomic O line fluxes."," Given the fact that the RMS variation in the $V$ and $R$ passbands is about 0.25 mag $^{-2}$ (see Table \ref{tab:dark}) ), this implies that the fluctuations seen in these passband are not completely accounted by the changes in the atomic O line fluxes."
+" The two lines appear to show a very weak correlation: the average ratio F(6300)/F(5577) is 0.64, but in a significant number of cases (~25%)) this ratio is larger than 1 (see also Sec. 7))."," The two lines appear to show a very weak correlation: the average ratio F(6300)/F(5577) is 0.64, but in a significant number of cases $\sim$ ) this ratio is larger than 1 (see also Sec. \ref{sec:corr}) )."
+" In those circumstances, the line becomes the most prominent nightglow atomic feature in the optical domain."," In those circumstances, the line becomes the most prominent nightglow atomic feature in the optical domain."
+" As far as the solar activity is concerned, it must be noticed that even though a dependence from the SFD is seen, it is less clear than in the case of the [OI]5577 (see Fig. 10,,"," As far as the solar activity is concerned, it must be noticed that even though a dependence from the $SFD$ is seen, it is less clear than in the case of the [OI]5577 (see Fig. \ref{fig:oi6300},"
+ lower panel)., lower panel).
+" In fact, even though the slopes are very similar, the linear correlation coefficient for [OI]6300 is 0.23, to be compared with 0.49 measured for the [OI]5577 line (see also Tab. 6))."," In fact, even though the slopes are very similar, the linear correlation coefficient for [OI]6300 is 0.23, to be compared with 0.49 measured for the [OI]5577 line (see also Tab. \ref{tab:lines}) )."
+" It is well known that, during the first hours of the night, the [OIJ6300 line shows a steady decline in brightness which, at tropical sites, is interrupted by abrupt emission bursts (see, for instance, Roach Gordon 1973))."," It is well known that, during the first hours of the night, the [OI]6300 line shows a steady decline in brightness which, at tropical sites, is interrupted by abrupt emission bursts (see, for instance, Roach Gordon \cite{roach}) )."
+ A plot of the line fluxes as a function of time elapsed after the evening twilight (see Fig. 12)), A plot of the line fluxes as a function of time elapsed after the evening twilight (see Fig. \ref{fig:decay}) )
+" clearly shows that these events occur during the whole night, and not only before midnight, as sometimes stated in the"," clearly shows that these events occur during the whole night, and not only before midnight, as sometimes stated in the"
+or each of these subsamples and. plot one against the other.,for each of these subsamples and plot one against the other.
+" While the variation of bias with formation time is smooth. he strongest. elfects clearly occur for the ""oldest haloes."," While the variation of bias with formation time is smooth, the strongest effects clearly occur for the “oldest” haloes."
+ sotice also that the variation in mean formation redshift is large. ranging from z=0.47 for the voungest to >=2.94 for the oldest.," Notice also that the variation in mean formation redshift is large, ranging from $z=0.47$ for the youngest to $z=2.94$ for the oldest."
+1054... Phe mean formation redshift or the population as a whole is 2=1.54., The mean formation redshift for the population as a whole is $z=1.54$.
+ 1n this Letter. we have used the very laree Alillennium Simulation (Springel et al 2005) to study how the clustering of dark haloes depends on mass ancl formation time.," In this Letter, we have used the very large Millennium Simulation (Springel et al 2005) to study how the clustering of dark haloes depends on mass and formation time."
+" Our results for the mean dependence. of bias on mass agree well with those of other workers. but for low mass haloes. Alo<Al,. we find a strong ancl unexpected: dependence on formation time."," Our results for the mean dependence of bias on mass agree well with those of other workers, but for low mass haloes, $M\leq M_*$, we find a strong and unexpected dependence on formation time."
+ Haloes of given mass that assembled at high. redshift are substantially more strongly clustered than haloes of the mass that assembled recently., Haloes of given mass that assembled at high redshift are substantially more strongly clustered than haloes of the mass that assembled recently.
+ This dillerence persists to large scale. and. indeed is larger on scales of order 20h.1 Mpc than on scales of order 2h 1Mpc.," This difference persists to large scale, and indeed is larger on scales of order $20h^{-1}$ Mpc than on scales of order $2h^{-1}$ Mpc."
+ AIrough there is unavoidably some arbitrariness in our definitions of halo mass anc formation time. experiments with alternative definitions of cach quantity produce bias variations which are typically of orderθά.," Although there is unavoidably some arbitrariness in our definitions of halo mass and formation time, experiments with alternative definitions of each quantity produce bias variations which are typically of order."
+ Our basic result thus appears robust., Our basic result thus appears robust.
+ This result. is unexpected because the sharp A-space filter used in most formulations of the exeursion set. mode! for structure formation (Bond et al 1991: Lacey Cole 1993) causes the random walks as a function of smoothing scale on which the mocel is based to have Markovian character., This result is unexpected because the sharp $k$ -space filter used in most formulations of the excursion set model for structure formation (Bond et al 1991; Lacey Cole 1993) causes the random walks as a function of smoothing scale on which the model is based to have Markovian character.
+ Vhe formation history of a halo is encoded in the random walk at higher mass resolution than that which defines the halo itself ancl is thus statistically independent of the environment. whieh is encoded in the random walk at lower resolution (White 1996).," The formation history of a halo is encoded in the random walk at higher mass resolution than that which defines the halo itself, and is thus statistically independent of the environment, which is encoded in the random walk at lower resolution (White 1996)."
+ “Phis independence of history and environment appeared confirmed by the simulation analyses of Lemson Ixaullmann (1999) ancl Percival et al.(2003)., This independence of history and environment appeared confirmed by the simulation analyses of Lemson Kauffmann (1999) and Percival et al.(2003).
+ These concentrated on relatively massive objects for which the effects. we find are small. probably undetectable with volumes of the size analysed.," These concentrated on relatively massive objects for which the effects we find are small, probably undetectable with volumes of the size analysed."
+" A formation time dependence was seen. however. in the mark correlation analysis of Sheth Tormen (2004). although these authors characterised the cllect as ""slight""."," A formation time dependence was seen, however, in the mark correlation analysis of Sheth Tormen (2004), although these authors characterised the effect as “slight”."
+ Independence of history anc environment for haloes of elven mass is assumed in many theoretical models for galaxy ormation., Independence of history and environment for haloes of given mass is assumed in many theoretical models for galaxy formation.
+ For example. Ixaulfmann. Nusser Steinmetz (1997). Benson et al (2000) ancl Wechsler et al (2001) populate haloes in dark matter simulations by using semi-analytic mocels applied to a Monte Carlo realisation of cach wlo’s history depending only on its mass.," For example, Kauffmann, Nusser Steinmetz (1997), Benson et al (2000) and Wechsler et al (2001) populate haloes in dark matter simulations by using semi-analytic models applied to a Monte Carlo realisation of each halo's history depending only on its mass."
+ Halo occupation distribution models such as those of Jing. Mo Borrncr 998). Peacock Smith (2000). Berlind et al (2003) or van den Bosch. Yang Alo (2003) ignore formation histories altogether and assume that the galaxy. population ofa halo depends on mass alone. independent (atfired mass) of the larger scale environment.," Halo occupation distribution models such as those of Jing, Mo Börrner (1998), Peacock Smith (2000), Berlind et al (2003) or van den Bosch, Yang Mo (2003) ignore formation histories altogether and assume that the galaxy population of a halo depends on mass alone, independent (at mass) of the larger scale environment."
+ In practice. haloes similar in mass to that of the Milky Way contain a substantial fraction of the galaxies in typical observational surveys.," In practice, haloes similar in mass to that of the Milky Way contain a substantial fraction of the galaxies in typical observational surveys."
+ Since it djs plausible that ealaxxy properties should. depend: significantly on the assembly history of their haloes. our results suggest that models which ignore the age dependence of halo clustering will incorrectly predict the large-scale distribution of galaxies.," Since it is plausible that galaxxy properties should depend significantly on the assembly history of their haloes, our results suggest that models which ignore the age dependence of halo clustering will incorrectly predict the large-scale distribution of galaxies."
+ The extent of the problem may depend on the specilic galaxy formation model considered., The extent of the problem may depend on the specific galaxy formation model considered.
+" We will estimate its size for a typical ""successful"" model in a future paper.", We will estimate its size for a typical “successful” model in a future paper.
+ Finally we note that. galaxy. formation models. which explicitly follow the assembly history of each halo should be immune to this problem., Finally we note that galaxy formation models which explicitly follow the assembly history of each halo should be immune to this problem.
+ Examples of such models. which eraft à semi-analvtic treatment of barvon physics onto a N-body. simulation can be found in Ixauffmann ct al (1999). Springe ct al (2001. 2005). Lelly et al (2003). Hatton et al (2003) and. Wang et al (2005).," Examples of such models which graft a semi-analytic treatment of baryon physics onto a N-body simulation can be found in Kauffmann et al (1999), Springel et al (2001, 2005), Helly et al (2003), Hatton et al (2003) and Kang et al (2005)."
+ Cosmological simulations which follow the harvonic physies. directly. should. also be safe. although it is not vet computationally feasible: to simulate galaxy formation reliably throughout a big enough volume for the elfects discussed here to dominate the errors.," Cosmological simulations which follow the baryonic physics directly should also be safe, although it is not yet computationally feasible to simulate galaxy formation reliably throughout a big enough volume for the effects discussed here to dominate the errors."
+ We are grateful to Adrian Jenkins anc Carlos Ετος for detailed comments on our paper., We are grateful to Adrian Jenkins and Carlos Frenk for detailed comments on our paper.
+ GL also. thanks Yipeng Jing and Darren Croton for useful discussions., GL also thanks Yipeng Jing and Darren Croton for useful discussions.
+ The simulation used in this paper was carried out as part of the programme of the Virgo Consortium on the Regatta supercomputer of the Computing Centre of the in Garching., The simulation used in this paper was carried out as part of the programme of the Virgo Consortium on the Regatta supercomputer of the Computing Centre of the Max-Planck-Society in Garching.
+"affected as much by uncertainties in the effective temperature, but more to uncertainties in the surface gravity through the dependence of the H bound-free opacity on electron pressure, and the iron abundance.","affected as much by uncertainties in the effective temperature, but more to uncertainties in the surface gravity through the dependence of the $^-$ bound-free opacity on electron pressure, and the iron abundance."
+ All the measured sulphur lines are weak and therefore not much affected by the microturbulence., All the measured sulphur lines are weak and therefore not much affected by the microturbulence.
+" According to ? the non-LTE corrections for the 1045 nm triplet are large and negative, i.e. the LTE abundances should be diminished in order to correct for the effects."," According to \cite{Takeda2005} the non-LTE corrections for the 1045 nm triplet are large and negative, i.e. the LTE abundances should be diminished in order to correct for the effects."
+" They have calculated non-LTE corrections for a grid of model atmospheres with Teg=4500,5000,5500,6000,6500,7000 K, logg-1.0,2.0,3.0,4.0,5.0, and [Fe/H]-4.0,—3.0,-2.0,21.0,—0.5,0.0, 0.5."," They have calculated non-LTE corrections for a grid of model atmospheres with $T_{\mathrm{eff}}=4500, 5000, 5500, 6000, 6500, 7000$ K, $\log g=1.0, 2.0, 3.0, 4.0, 5.0$, and $\left[\mathrm{Fe}/\mathrm{H}\right]=-4.0, -3.0, -2.0, -1.0, -0.5, 0.0, +0.5$ ."
+ Their calculations use hydrogen collision cross sections based on the classical formula by ?? in the version of ? ie. Sy=1 is assumed.," Their calculations use hydrogen collision cross sections based on the classical formula by \citet{Drawin1968,Drawin1969} in the version of \citet{Steenbock1984} i.e. $S_H=1$ is assumed."
+" The work of ? is, to our knowledge, the only available non-LTE corrections for the 1045 nm triplet."," The work of \cite{Takeda2005} is, to our knowledge, the only available non-LTE corrections for the 1045 nm triplet."
+" ? has, however, also performed non-LTE calculations for sulphur, but for the multiplet around 675 nm, the 869 nm doublet, and the 923 nm triplet."," \cite{Korotin2009} has, however, also performed non-LTE calculations for sulphur, but for the multiplet around 675 nm, the 869 nm doublet, and the 923 nm triplet."
+" The overlapping parts of the two works agree well for lower-temperature stars, but differ for temperatures above 6000 K. The non-LTE correction used (?) for our observed stars is shown in Table B]."," The overlapping parts of the two works agree well for lower-temperature stars, but differ for temperatures above 6000 K. The non-LTE correction used \citep{Takeda2005} for our observed stars is shown in Table \ref{tab:abund}."
+" Note that we have not attempted to interpolate in the grid, but chosen the corrections from the models best applicable to the stars in this study."," Note that we have not attempted to interpolate in the grid, but chosen the corrections from the models best applicable to the stars in this study."
+ We note the slight inconsistency of using spherical model photospheres for our giants while applying non-LTE corrections calculated in plane parallel geometry., We note the slight inconsistency of using spherical model photospheres for our giants while applying non-LTE corrections calculated in plane parallel geometry.
+ Further investigations on the effects of non-LTE corrections due to spherical geometry would be useful., Further investigations on the effects of non-LTE corrections due to spherical geometry would be useful.
+" Regarding the 1082 nm [Si] line, being an inter-combination resonance line, LTE is expected to be a good approximation, because the populations of the energy levels are determined by collisions rather than radiative transitions."," Regarding the 1082 nm $[\ion{S}{i}] $ line, being an inter-combination resonance line, LTE is expected to be a good approximation, because the populations of the energy levels are determined by collisions rather than radiative transitions."
+"outside a radius of 1"" are simply allocated their expected poisson error during the model fitting process.",outside a radius of $1\asec$ are simply allocated their expected poisson error during the model fitting process.
+ However. Vig 2 also demonstrates that this system. would: grosslv underestimate. the actual error introduced. inside-- 1H bv he severe undersampling of the WEPC2 PSE.," However, Fig \ref{errprof} also demonstrates that this system would grossly underestimate the actual error introduced inside $\simeq1\asec$ by the severe undersampling of the WFPC2 PSF."
+" Therefore. inside a radius of 1"" a different svstem is used to allocate he error weighting."," Therefore, inside a radius of $1\asec$ a different system is used to allocate the error weighting."
+ A series of ten pixel-wice circular annuli centred. on the quasar are constructed. and the variance (o7) of the distribution of the pixels falling within cach annuli calculated.," A series of ten pixel-wide circular annuli centred on the quasar are constructed, and the variance $\sigma^{2}$ ) of the distribution of the pixels falling within each annuli calculated."
+ All of the pixels falling within a particular annulus are then allocated. the annulus variance as their error weighting., All of the pixels falling within a particular annulus are then allocated the annulus variance as their error weighting.
+ This procedure is still justified even when the underlving host galaxy has a clear elipticity since. even at resolution. the central 17 of the| (quasar images are dominated by the circular svimetry of he PSE.," This procedure is still justified even when the underlying host galaxy has a clear ellipticity since, even at resolution, the central $1\asec$ of the quasar images are dominated by the circular symmetry of the PSF."
+" Due to the success of the re-sampling technique disc""ussed above. ibi was possible to assign a poisson error welgMing to the central pixel without biasing the fitting procedure."," Due to the success of the re-sampling technique discussed above, it was possible to assign a poisson error weighting to the central pixel without biasing the fitting procedure."
+ This complicated process has proven to be successful. wih the vast majority of the 33 objects modelled having minimum x solutions lving in the v=V2» region expected. fcX correct weighting Dunlop et al.," This complicated process has proven to be successful, with the vast majority of the 33 objects modelled having minimum $\chi^{2}$ solutions lying in the $\nu\pm\sqrt{2\nu}$ region expected for correct weighting (Dunlop et al."
+ 2000)., 2000).
+ Confidence in the weighting svstem is further strengthened by examination of the pixel maps of reduced. 47 values produced. for the best-fitting mocels., Confidence in the weighting system is further strengthened by examination of the pixel maps of reduced $\chi^{2}$ values produced for the best-fitting models.
+ Shown in Fig 3. a typical A7 map from the best-fit to the mraclio-loud quasar 1217)23., Shown in Fig \ref{chimap} a typical $\chi^{2}$ map from the best-fit to the radio-loud quasar 1217+023.
+ Ht is clear from this figure that 10. desired uniform spread of X7 values has been achieved. with no one area of the image dominating the fit.," It is clear from this figure that the desired uniform spread of $\chi^{2}$ values has been achieved, with no one area of the image dominating the fit."
+ The main advantage of the two-dimensional mocelling technique being adopted here is that dt allows each individual pixel in the image frame to x: included. in the model fitting as a degree of freedom., The main advantage of the two-dimensional modelling technique being adopted here is that it allows each individual pixel in the image frame to be included in the model fitting as a degree of freedom.
+ However. due to the high angular resolution of our data. a tvvical model fit to à radius of 12” will include some 45000 WE£2 pixels (~5700 IRCAAI pixels) in cach 47 evaluation.," However, due to the high angular resolution of our data, a typical model fit to a radius of $12\asec$ will include some 45000 WFPC2 pixels $\sim5700$ IRCAM3 pixels) in each $\chi^{2}$ evaluation."
+ E10 adopted solution was to utilise the downhill simplex. methoc (Press et al., The adopted solution was to utilise the downhill simplex method (Press et al.
+ 1989) to locate the global minimum., 1989) to locate the global minimum.
+ Altrough significantly. slower than the conjugate gradient metrock used in. Taylor et al. (, Although significantly slower than the conjugate gradient method used in Taylor et al. (
+1996). downhill simplex minimizaion was chosen for,"1996), downhill simplex minimization was chosen for"
+Some of the iron Ίνα Lue observations - particularly in the realii of variability - have been surprising in the coutext of the standL disk-line model.,Some of the iron $\alpha$ line observations - particularly in the realm of variability - have been surprising in the context of the standard disk-line model.
+ Depeudiug ou ones viewpoint. this may be interpreted as a furtlier demonstration of their diagnostic power (see above) or as being problematic fx the disk line model (e.9.. Suleutic. Marziani Calvani 199842).," Depending on ones viewpoint, this may be interpreted as a further demonstration of their diagnostic power (see above) or as being problematic for the disk line model (e.g., Sulentic, Marziani Calvani 1998a)."
+ Alternatives to the relativistic disk have been suggested and we review those briefly here., Alternatives to the relativistic disk have been suggested and we review those briefly here.
+ Maux of thexdastes have been discussed by Fabi et al. (, Many of these issues have been discussed by Fabian et al. (
+1995).,1995).
+ They include:geometries: the original papey predicting the iron Ίνα lines. Caulbert Rees (1988) sugeestedOO that they may come from optically thick clouds. sheets or filaments. rather than the disk.," They include: in the original paper predicting the iron $\alpha$ lines, Guilbert Rees (1988) suggested that they may come from optically thick clouds, sheets or filaments, rather than the disk."
+ Such material can produce the observed. line streugths. as long as the covering fraction is high. but less than unity to avoid obscuring the contiuuun (c.g. Bond Matsuoka 1993: Naudra George 1991).," Such material can produce the observed line strengths, as long as the covering fraction is high, but less than unity to avoid obscuring the continuum (e.g. Bond Matsuoka 1993; Nandra George 1994)."
+ Detailed caleulatious of the profiles iu a pseucdo-spherical geometry have not beeu carried out aud it is therefore difficult to make a definite statement as to whether the ο MUNcan be reproduced., Detailed calculations of the profiles in a pseudo-spherical geometry have not been carried out and it is therefore difficult to make a definite statement as to whether the observations can be reproduced.
+ It does. however. seen uulikelv.," It does, however, seem unlikely."
+ Iu. MCCC-6-30-15 and NGC and35 (Fig., In MCG-6-30-15 and NGC 3516 (Fig.
+" 1). there is very little evidence for emission bluc-wiud of the rest (rev, is hard to euvisage a non-disk οἳcolctry in which the blue emission is suppressed."," 1), there is very little evidence for emission blue-ward of the rest energy, and it is hard to envisage a non-disk geometry in which the blue emission is suppressed."
+ The details depend on the ecometry. kinenaaties aud self-covering of the cloud system but. for example rotating blobs are alinost certainly excluded. given wat the high velocities would cause Doppler boosting of the blue wine.," The details depend on the geometry, kinematics and self-covering of the cloud system but, for example rotating blobs are almost certainly excluded, given that the high velocities would cause Doppler boosting of the blue wing."
+ Inflowine or outflowing cloud distributions are another possibility. but once again it secs ard to suppress the cussion from material moving towards us without a special ecouctry.," Inflowing or outflowing cloud distributions are another possibility, but once again it seems hard to suppress the emission from material moving towards us without a special geometry."
+ As moeutioued above. Zaziarski et al. (," As mentioned above, Zdziarski et al. ("
+1996) have suggested hat an appareuth-broad line im NGC 1151 can be accounted for by a complex continuum - in particular complex aML,1996) have suggested that an apparently-broad line in NGC 4151 can be accounted for by a complex continuum - in particular complex absorption.
+ Many AGN have material in their ines-of-xieht. such as “wari absorbers” (6.8. Reynolds 1997: Georee et al.," Many AGN have material in their lines-of-sight, such as “warm absorbers” (e.g. Reynolds 1997; George et al."
+ 1998) mt typically these are not thought to affect the spectrum above about 3 keV. As in all other svavebands. a great challenge in determining the hue streneths aud xofiles is deteziniue the wuderlvingo continua.," 1998) but typically these are not thought to affect the spectrum above about 3 keV. As in all other wavebands, a great challenge in determining the line strengths and profiles is determining the underlying continuum."
+ Although different continuumdeconvolution methods tend to produce similar results. uukuown complexities add additional uncertainty to the profile determination.," Although different continuum-deconvolution methods tend to produce similar results, unknown complexities add additional uncertainty to the profile determination."
+ Winematics aud gravitatk rare not the only mechanisis by which lines can be broadened., Kinematics and gravitation are not the only mechanisms by which lines can be broadened.
+ A particular alternative sueecstionOO las been that Compton scattering can broaden the line (6.g.. Misra Whembavi 1998).," A particular alternative suggestion has been that Compton scattering can broaden the line (e.g., Misra Khembavi 1998)."
+ Fabian et al. (, Fabian et al. (
+1995) argued against such a model. as the Comptonizine mediuu must be finely-tuned to ecuerate the observed profile.,"1995) argued against such a model, as the Comptonizing medium must be finely-tuned to generate the observed profile."
+ Iu particulary a very compact. lieh optical depth (7~ 5). cool (~0.2 keV) mediani ds required.," In particular a very compact, high optical depth $\tau \sim 5$ ), cool $\sim
+0.2$ keV) medium is required."
+ Such a medium would dowu-scatter the N-vav coutinuun also. proclucing a spectral break at ~20 keV which is not observed.," Such a medium would down-scatter the X-ray continuum also, producing a spectral break at $\sim 20$ keV which is not observed."
+ Combinations of numerous ines can appear broad when observed at low resolution., Combinations of numerous lines can appear broad when observed at low resolution.
+ Differently-ionized. species of iron-k. for example. result in enission from 6.1-6.9 keV. This cannot explain the observed profiles. as the broad emissiou occurs below 6.1 keV. No abundant element xoduces ine emission in the 1-0 keV," Differently-ionized species of iron-K, for example, result in emission from 6.4-6.9 keV. This cannot explain the observed profiles, as the broad emission occurs below 6.4 keV. No abundant element produces line emission in the 4-6 keV"
+lower-limit for the magnetic field.,lower-limit for the magnetic field.
+" In the Thomson regime. this is given by writing ys= y,/107. T,;=T,/40000 K and Ryo=R,/10Rs the temperature and radius of the companion star. and do,= d/O.] AU the orbital separation."," In the Thomson regime, this is given by writing $\gamma_3=\gamma_e/10^3$ , $T_{\star,4}=T_{\star}/40~000$ K and $R_{\star,10}=R_{\star}/{10~R_{\odot}}$ the temperature and radius of the companion star, and $d_{0.1}=d/0.1$ AU the orbital separation."
+" Using the approximate formula for Compton energy losses (?).. the same condition in the extreme Klein-Nishina regime holds if If the Larmor radius is compared with the Compton mean free path given by Uj.~πόση. where n, is the stellar photon density and ο. the Compton cross section. the condition on the magnetic field is more restrictive."," Using the approximate formula for Compton energy losses \citep{1970RvMP...42..237B}, the same condition in the extreme Klein-Nishina regime holds if If the Larmor radius is compared with the Compton mean free path given by $\lambda_{ic}\sim 1/n_{\star}\sigma_{ic}$, where $n_{\star}$ is the stellar photon density and $\sigma_{ic}$ the Compton cross section, the condition on the magnetic field is more restrictive."
+" In the Thomson regime. the electron loses only a small fraction of its total energy per interaction. hence l4,>Aj."," In the Thomson regime, the electron loses only a small fraction of its total energy per interaction, hence $\lambda_{cool}>\lambda_{ic}$."
+ In the Klein-Nishina regime. most of the electron energy is lost in a single scattering and luu3Ai.," In the Klein-Nishina regime, most of the electron energy is lost in a single scattering and $\lambda_{cool}\approx\lambda_{ic}$."
+ Because the cascade occurs mostly in the Klein-Nishina regime in gamma-ray binaries. both conditions lead approximatively to the same lower limit for the ambient magnetic field.," Because the cascade occurs mostly in the Klein-Nishina regime in gamma-ray binaries, both conditions lead approximatively to the same lower limit for the ambient magnetic field."
+ In addition to this condition. pairs are assumed to be isotropized at their creation site for simplicity.," In addition to this condition, pairs are assumed to be isotropized at their creation site for simplicity."
+ Pairs will be randomized if the ambient magnetic field is disorganized., Pairs will be randomized if the ambient magnetic field is disorganized.
+ Isotropization of pairs in the cascade will also occur due to pitch angle scattering if the magnetic turbulence timescale Is smaller than the energy loss timescale 1f itis on the order of the Larmor timescale)., Isotropization of pairs in the cascade will also occur due to pitch angle scattering if the magnetic turbulence timescale is smaller than the energy loss timescale if it is on the order of the Larmor timescale).
+ For lower magnetic field intensity Cantsotropic’ domain in Fig. 1).," For lower magnetic field intensity (`anisotropic' domain in Fig. \ref{domain}) ),"
+ the cascade remains three-dimensional but then pairs cannot be considered as locally isotropized., the cascade remains three-dimensional but then pairs cannot be considered as locally isotropized.
+ In this case. the trajectories of the particles should be properly computed as in ?..," In this case, the trajectories of the particles should be properly computed as in \citealt{2005MNRAS.356..711S}."
+ For B€1075 G. the cascade is one-dimensional (?)..," For $B\la 10^{-8}~$ G, the cascade is one-dimensional \citep{2009A&A...507.1217C}."
+ If the magnetic field is too strong. pairs locally isotropize but cool down via synchrotron radiation rather than by inverse Compton scattering.," If the magnetic field is too strong, pairs locally isotropize but cool down via synchrotron radiation rather than by inverse Compton scattering."
+ Most of the energy is then emitted in X-rays and soft gamma rays. below the threshold energy for par production.," Most of the energy is then emitted in X-rays and soft gamma rays, below the threshold energy for pair production."
+ The cascade is quenched as soon as the first generation of pairs is produced., The cascade is quenched as soon as the first generation of pairs is produced.
+ This condition gives an upper-limit for the magnetic field., This condition gives an upper-limit for the magnetic field.
+" Synchrotron losses are smaller than inverse Compton losses Ey,<E; for in the Thomson regime and for the deep Klein-Nishina regime (?).. ", Synchrotron losses are smaller than inverse Compton losses $\dot{E}_{syn}<\dot{E}_{ic}$ for in the Thomson regime and for in the deep Klein-Nishina regime \citep{1970RvMP...42..237B}. .
+It can be noticed that the most relevant upper-limit for the magnetic field strength ts given by the Thomson formula in Eq. (3)).," It can be noticed that the most relevant upper-limit for the magnetic field strength is given by the Thomson formula in Eq. \ref{bmax_th}) ),"
+" since high-energy particles (E,21 GeV) with Bey<B«By can cool down and get into the cascade domain.", since high-energy particles $E_e\ga 1$ GeV) with $B_{KN}<B<B_T$ can cool down and get into the cascade domain.
+" Figure | shows the complete domain where a three-dimensional ""isotropic? cascade can be initiated in5039.. combining the lower and upper-limit for B."," Figure \ref{domain} shows the complete domain where a three-dimensional `isotropic' cascade can be initiated in, combining the lower and upper-limit for $B$."
+ This domain encompasses plausible values for the ambient magnetic. field in the system., This domain encompasses plausible values for the ambient magnetic field in the system.
+" It is worthwhile to note that for very high-energy electrons E,245do)Bo, TeV. where By,=B/O.1 G. the Larmor radius becomes greater than the binary separation in (Fig. 1))."," It is worthwhile to note that for very high-energy electrons $E_e\ga 45~d_{0.1}B_{0.1}$ TeV, where $B_{0.1}=B/0.1~$ G, the Larmor radius becomes greater than the binary separation in (Fig. \ref{domain}) )."
+ In this case. the local magnetic confinement approximation of particles is not appropriate anymore.," In this case, the local magnetic confinement approximation of particles is not appropriate anymore."
+ This ts unlikely to happen in uf the VHE emission has a leptonic origin since HESS observations shows an energy cut-off for photons at «10 Τεν. Contrary to the one-dimensional case. three-dimensional pair cascading cannot be explicitly computed.," This is unlikely to happen in if the VHE emission has a leptonic origin since HESS observations shows an energy cut-off for photons at $\approx 10$ TeV. Contrary to the one-dimensional case, three-dimensional pair cascading cannot be explicitly computed."
+ Nevertheless. it is possible to decompose the cascade into successive generations of particles.," Nevertheless, it is possible to decompose the cascade into successive generations of particles."
+ Two different approaches are presented below. one based on semi-analytical calculations and the other on a Monte Carlo code.," Two different approaches are presented below, one based on semi-analytical calculations and the other on a Monte Carlo code."
+ In both models. the primary source of gamma rays is point-like and comeident with the compact object position as it is depicted in Fig. 2..," In both models, the primary source of gamma rays is point-like and coincident with the compact object position as it is depicted in Fig. \ref{geo}."
+ The origin and the angular dependence of the primary gamma-ray flux are not specified at this stage., The origin and the angular dependence of the primary gamma-ray flux are not specified at this stage.
+ These methods are general and could be applied to any other astrophysical context involving 3D pair cascading., These methods are general and could be applied to any other astrophysical context involving 3D pair cascading.
+ A beam of primary gamma rayspropagating in the direction defined by the spherical angles € and @ (see Fig. 2)).," A beam of primary gamma rayspropagating in the direction defined by the spherical angles $\theta$ and $\phi$ (see Fig. \ref{geo}) ),"
+ produces at adistance + to the primary source the first generation of pairs., produces at adistance $r$ to the primary source the first generation of pairs.
+ In the point-like and mono-energetie star approximation.," In the point-like and mono-energetic star approximation,"
+Alone our nearest ealactic neighbors. the Maeellauic Clouds. with their relative proximity. low lue-ofsight extinction. lower mean metalicitv. and large nuinbers of stellar populations. can provide us with vital information ὲut our own Milkv Wav aud galactic evolution iu eeneral,"Among our nearest galactic neighbors, the Magellanic Clouds, with their relative proximity, low line-of-sight extinction, lower mean metallicity, and large numbers of stellar populations, can provide us with vital information about our own Milky Way and galactic evolution in general."
+ A uultitude of studies at a large variety of wavoeleugtis have been conducted on the Clouds for decades., A multitude of studies at a large variety of wavelengths have been conducted on the Clouds for decades.
+ Larec-scale optical strveys. such as the Magellanic Couds Survey (Zaritsky. Warvis. Thompson 1997). are coupiling results on a large nuunber of stellar objects aud uebulae.," Large-scale optical surveys, such as the Magellanic Clouds Survey (Zaritsky, Harris, Thompson 1997), are compiling results on a large number of stellar objects and nebulae."
+ To obtain a view particularly of the post-iuiaiu-secuence stellar populations in these galaxies oue needs to furn to the infrared (IR)., To obtain a view particularly of the post-main-sequence stellar populations in these galaxies one needs to turn to the infrared (IR).
+ Two recent infrared survevs have xovided a wealth of new data ou the Large Magellanic €oud (LMC) in the wavelength regime from 1.2 to 25 ju. At near-IR wavelengths. the Two Micron All Sky Survey (2\MLASS) has detected nearly 2 uullion stars in the 107«1t lg area also surveved by the Midcourse Space Expernucut (AISN) iid-IB. survey.," Two recent infrared surveys have provided a wealth of new data on the Large Magellanic Cloud (LMC) in the wavelength regime from 1.2 to 25 $\mu$ m. At near-IR wavelengths, the Two Micron All Sky Survey (2MASS) has detected nearly 2 million stars in the $10%
+\arcdeg\times10\arcde g area also surveyed by the Midcourse Space Experiment (MSX) mid-IR survey."
+ AISX provies an Auprovenieunt over previous nuappiug of the LMC in the mid-IR by the US/Netherlauds/Ul IR. Astronomical Saelite (IRAS)., MSX provides an improvement over previous mapping of the LMC in the mid-IR by the US/Netherlands/UK IR Astronomical Satellite (IRAS).
+ The smaller detector footprint of fie MISN detector arrays. as compared to IRAS. avoids he confusion problems eucountered by the IRAS survey i sole areas of the LMC. aud allows a more precise detenuinatio1 of source position. Which is of ercat inportance Ἡ“HOLL atempting to compare these mid-IR observations with the 2\LTASS near-IR observations.," The smaller detector footprint of the MSX detector arrays, as compared to IRAS, avoids the confusion problems encountered by the IRAS survey in some areas of the LMC, and allows a more precise determination of source position, which is of great importance when attempting to compare these mid-IR observations with the 2MASS near-IR observations."
+ The AISX A baud 1 is also more sensitive thau the IRAS 12 jn1|G.8xuxd [7 25 pna]nu]. so therefore we can catalog LsΙ sources to a limiting uaenitude —7.5.," The MSX A band [6.8 – 10.8 $\mu $ m] is also more sensitive than the IRAS 12 $\mu $ m band [7 – 15 $\mu $ m], so therefore we can catalog 1806 sources to a limiting magnitude $\sim $ 7.5."
+ The utility of mid-IR observatious in icdentifving post-uazin-sequenuce stars in the LAIC was demoustrated by. IRAS. which.wi ha detection limit of magnitude ~6. detected several lance extreme evolved stars (such as ΟΠΤΗ stars) in the LAIC.," The utility of mid-IR observations in identifying post-main-sequence stars in the LMC was demonstrated by IRAS, which,with a detection limit of magnitude $\sim 6$, detected several hundred extreme evolved stars (such as OH/IR stars) in the LMC."
+ At the distance modulus of the LMC. ouly the extremely red objects at the brighter οid of the iuid-IR ΠΟ]Τ1de scale could be detected by IRAS.," At the distance modulus of the LMC, only the extremely red objects at the brighter end of the mid-IR magnitude scale could be detected by IRAS."
+ By extending the LMC iiid-IR 3rvev nearly two magitudes fainter we expect fo see some of the less extreme members of the LMC eiant brauch »opulations aud other objecs., By extending the LMC mid-IR survey nearly two magnitudes fainter we expect to see some of the less extreme members of the LMC giant branch populations and other objects.
+ 2MASSobservations. extending to J= 15.," 2MASSobservations, extending to $%
+J=15."
+5 nage a a signalto-noise ratio (SNR) of 10. have detect «Land vieled information on firs-ascent red giant branch. :isvinptotic eiut branch (AGB). and obscured AGB »opulaions of he LAIC (Nikolaev Weinberg 2000).," 8 mag at a signal-to-noise ratio (SNR) of 10, have detected and yielded information on first-ascent red giant branch, asymptotic giant branch (AGB), and obscured AGB populations of the LMC (Nikolaev Weinberg 2000)."
+ These red populatious have also been iuveΠοιος aud analyzed using he DENIS point source catalogs by Cioni et al. (, These red populations have also been inventoried and analyzed using the DENIS point source catalogs by Cioni et al. (
+2XWOO) and otos papers bv the DENTS team.,2000) and other papers by the DENIS team.
+ By combiniig tie near-IR and nud-IR data. we ereatly expand the color bascli1ο and acquire additional insight iuto the ACB aud ot101) red stars m the LAIC. as well as discoveriic objects with unsual IR excesses.," By combining the near-IR and mid-IR data, we greatly expand the color baseline and acquire additional insight into the AGB and other red stars in the LMC, as well as discovering objects with unusual IR excesses."
+ Tu this paper we present the resits of our ποσο the MSX aud 2MASS LAC datasets axd discuss the nature of the various IR-bright sources in that ealaxy. manv of which are previously unideuti&ed.," In this paper we present the results of our merging of the MSX and 2MASS LMC datasets and discuss the nature of the various IR-bright sources in that galaxy, many of which are previously unidentified."
+ TIus study provides a foundation for future strucies of the LAC by the SIRTF aud missions., This study provides a foundation for future studies of the LMC by the and missions.
+This basic characterization demonstrates the potential of using the X-ray tomography to greatly advance our understanding of the global hot gas in and around the Galaxy.,This basic characterization demonstrates the potential of using the X-ray tomography to greatly advance our understanding of the global hot gas in and around the Galaxy.
+ 'The X-ray absorption line spectroscopy can also be used to probe the hot CGM around other galaxies., The X-ray absorption line spectroscopy can also be used to probe the hot CGM around other galaxies.
+ Yao et al., Yao et al.
+ have recently studied the hot CGM along the sight lines toward luminous AGNs (24).., have recently studied the hot CGM along the sight lines toward luminous AGNs \cite{yao10}.
+" As an example, Fig."," As an example, Fig."
+ 2 shows a ggrating spectrum of PKS 2155-304., 2 shows a grating spectrum of PKS 2155-304.
+" No absorption line is apparent at the redshifts of the three groups of galaxies with the impact distances smaller than 500 kpc of the sight line (in particular, the group at the highest redshift contains two galaxies brighter than L,.)."," No absorption line is apparent at the redshifts of the three groups of galaxies with the impact distances smaller than 500 kpc of the sight line (in particular, the group at the highest redshift contains two galaxies brighter than $L_*$ )."
+ The same is true when the spectrum is folded according to the redshifts of the individual groups., The same is true when the spectrum is folded according to the redshifts of the individual groups.
+" To increase the counting statistics further, Yao et al."," To increase the counting statistics further, Yao et al."
+ have stacked spectra from the ggrating observations of eight AGNs with such intervening galaxies or groups., have stacked spectra from the grating observations of eight AGNs with such intervening galaxies or groups.
+ No significant absorption is detected for any of these individual systems or in the final stacked spectrum with a total equivalent exposure of about 10 Ms!, No significant absorption is detected for any of these individual systems or in the final stacked spectrum with a total equivalent exposure of about 10 Ms!
+ Upper limits to the mean column densities of various ion species per galaxy, Upper limits to the mean column densities of various ion species per galaxy
+" where ©, is the integrated number density of galaxies. σ, is the characteristic velocity dispersion. a is (he low-velocity power-law index. and 3 is the high-velocity exponential eut- index.","] where $\phi_*$ is the integrated number density of galaxies, $\sigma_*$ is the characteristic velocity dispersion, $\alpha$ is the low-velocity power-law index, and $\beta$ is the high-velocity exponential cut-off index."
+ Shethetal.(2003) (the nmunber density being updated by Mitchelletal.2005)) found from the SDSS DRI for the earlv-tvpe galaxy. population where / is the IIubble constant in units of LOO kms+ 7., \citet{She03} (the number density being updated by \citealt{Mit05}) ) found from the SDSS DR1 for the early-type galaxy population where $h$ is the Hubble constant in units of 100 $\kms$ $^{-3}$.
+ This was the first direct measurement of the VDF of eurlv-tvpe galaxies., This was the first direct measurement of the VDF of early-type galaxies.
+ Just recently. Choietal.(2007) have found from the much larger SDSS DR5 The above VDFs are intrinsic finctions derived. taking into account the correlated measurement errors of σ (Shethetal.2003:Choi 2007).," Just recently, \citet{Cho06} have found from the much larger SDSS DR5 The above VDFs are intrinsic functions derived taking into account the correlated measurement errors of $\sigma$ \citealt{She03,Cho06}) )."
+ The DRI and DR5 VDESs are shown in Figure 1((a)., The DR1 and DR5 VDFs are shown in Figure \ref{fig:Om}( (a).
+ The DR5 VDF is clearly quite different from the DRI VDF both in (he number density aud the shape for the lower part of σ., The DR5 VDF is clearly quite different from the DR1 VDF both in the number density and the shape for the lower part of $\sigma$.
+ This is in large part due to the improved galaxy classification scheme of Park&Choi(2005).. who makes use of à SDSS u-r color versus g-i color gradient space.," This is in large part due to the improved galaxy classification scheme of \citet{PC05}, who makes use of a SDSS $u$ $r$ color versus $g$ $i$ color gradient space."
+ While earlv-tvpe galaxies dominate strong lensing. late-tvpe galaxies cannot be neglected.," While early-type galaxies dominate strong lensing, late-type galaxies cannot be neglected."
+ Among the racdio-selected lenses with known galaxy tvpes from the CLASS and the PANELS (see Table 1 of Chae 2005)). about are late-tvpes.," Among the radio-selected lenses with known galaxy types from the CLASS and the PANELS (see Table 1 of \citealt{Cha05}) ), about are late-types."
+ For the Iate-tvpe galaxy population. the direct measurement of the VDF is complicated by Che significant rotations of the disks.," For the late-type galaxy population, the direct measurement of the VDF is complicated by the significant rotations of the disks."
+ We proceed as follows., We proceed as follows.
+ We adopt the Iuminosity hunction (LE) of late-tvpe galaxies [rom the SDSS DI5 (Choiοἱal. 2007))., We adopt the luminosity function (LF) of late-type galaxies from the SDSS DR5 \citealt{Cho06}) ).
+ We turn the LF into a circular velocity ποΊο using a relation (Tully&Pierce 2000))., We turn the LF into a circular velocity function using a Tully-Fisher relation \citealt{TP00}) ).
+ Finally. we tiun the circular velocity [uncetion into a VDF assuming that (he circular velocity is proportional to the inner velocity dispersion as would be the case in an isothermal model.," Finally, we turn the circular velocity function into a VDF assuming that the circular velocity is proportional to the inner velocity dispersion as would be the case in an isothermal model."
+ In principle. we can estimate all (he parameters of," In principle, we can estimate all the parameters of"
+introduce a slow decay of the field to produce polar fields which reverse each cycle.,introduce a slow decay of the field to produce polar fields which reverse each cycle.
+" The reconstructions considerably extend the basis for correlations studies, such as the relation between the polar field amplitude during activity minima and the strengths of the preceding and subsequent cycles, with implications for dynamo models."," The reconstructions considerably extend the basis for correlations studies, such as the relation between the polar field amplitude during activity minima and the strengths of the preceding and subsequent cycles, with implications for dynamo models."
+ Introducing a possibly ‘lost’ cycle at the end of the 18th century leads to a shift of the open flux minimum during the Dalton minimum which is incompatible with the '0Be , Introducing a possibly `lost' cycle at the end of the 18th century leads to a shift of the open flux minimum during the Dalton minimum which is incompatible with the $^{10}$ 
+associated with SNR 0353.6-0.7.,associated with SNR G353.6-0.7.
+ Typically found. iu vouug SNRs. CCOs are characterized by. blackbodyike soft N-rav cluission with temperatures of 0.2-0.5 τον) {Pavlovctal.2001).," Typically found in young SNRs, CCOs are characterized by blackbody-like soft X-ray emission with temperatures of 0.2-0.5 keV \citep{pav04}."
+. Tt is sugeested that CCOs are neutron stars born iu supernova explosions with xoperties different from those of classical rotation-vowered pulsars (Gotthelfetal.2005:Li2007).," It is suggested that CCOs are neutron stars born in supernova explosions with properties different from those of classical rotation-powered pulsars \citep{got05, li07}."
+".. NMMS J17320:3-311518 shows a blackbody temperature (0.5 seV) hügher than most CCOs,", XMMS J173203-344518 shows a blackbody temperature $\sim$ 0.5 keV) higher than most CCOs.
+ We think that it is possibly a maenetar (0. ANPs aud SCRs). because maguetars ive soft spectra. a typical blackbody. temperature of —0.5 keV and a typical luminosity of 1001°° ere &! (Mereehetti2008).. also see a inaguetarC," We think that it is possibly a magnetar (i.e. AXPs and SGRs), because magnetars have soft spectra, a typical blackbody temperature of $\sim$ 0.5 keV and a typical luminosity of $^{34-36}$ erg $^{-1}$ \citep{mer08}, also see a magnetar."
+"atalogue’, Given a distance of ~ 3.2 kpc. i£ NMMS J173203-311518 is associated with the SNR. it will have a 1-10. keV tuninosity of ~ 10°! ere +."," Given a distance of $\sim$ 3.2 kpc, if XMMS J173203-344518 is associated with the SNR, it will have a 1-10 keV luminosity of $\sim$ $^{34}$ erg $^{-1}$."
+ This perhaps ucaus it is like ANPs more than SCRs. because SGRs usually wave DIuuinosities higher than ANPs. and this source so ar shows no repeated soft οταν enüssion.," This perhaps means it is like AXPs more than SGRs, because SGRs usually have luminosities higher than AXPs, and this source so far shows no repeated soft $\gamma$ -ray emission."
+ We examine he aud Suzuka light curves of NMMS J173203-311518 and fid uo evidence for pulsations at frequencies between 0.05-1 Uz. a range typical of naenetars.," We examine the and Suzuka light curves of XMMS J173203-344518 and find no evidence for pulsations at frequencies between 0.05-1 Hz, a range typical of magnetars."
+ Although maguctars offen show a hard rav (tens of keV) tail that can be fitted bv a shallow over Lav model. no counterpart of NAIAIS J173203-311518 is found in the INTECGRAL/ISCGRI map (in the cherev range of 20-10 keV) aud INTEGRAL/JEMN nap (3-15 keV).," Although magnetars often show a hard X-ray (tens of keV) tail that can be fitted by a shallow power law model, no counterpart of XMMS J173203-344518 is found in the INTEGRAL/ISGRI map (in the energy range of 20-40 keV) and INTEGRAL/JEMX map (3-15 keV)."
+ Iu addition. we notice that there is a ittle higher cohuun density in the direction of NMMS J17320:3-3115185 than from other regions of SNR O353.6-(see Table 1).," In addition, we notice that there is a little higher column density in the direction of XMMS J173203-344518 than from other regions of SNR G353.6-0.7(see Table 1)."
+ This may be caused by a local dense cloud., This may be caused by a local dense cloud.
+ Future N-rayv observations are needed to test the uaenetar scenario for NAIAIS J173203-311518., Future X-ray observations are needed to test the magnetar scenario for XMMS J173203-344518.
+ Based on a statistics of 5 vouug SNRs. the typical photon index of ~ 2.1 - 3.7 for non-thermal SNRs las been sugeested (Bambaetal.2005).," Based on a statistics of 5 young SNRs, the typical photon index of $\sim$ 2.1 - 3.7 for non-thermal SNRs has been suggested \citep{bam05}."
+. It is interesting that the hard ταν cussion of the rine iu the TESS J1731-217/SNR. 0353.6-0.7 has similar photon iudex of ~ 2.3., It is interesting that the hard X-ray emission of the ring in the HESS J1731-347/SNR G353.6-0.7 has similar photon index of $\sim$ 2.2.
+ But the SNR is old due to its exteuded (~ S0 απο in diuueter) aud faint radio cnussion CIUS) so it should have a different higl-energw mechaisin., But the SNR is old due to its extended $\sim$ 30 arcmin in diameter) and faint radio emission (T08) so it should have a different high-energy mechanism.
+ Old SNRs have quite different characteristics for the expected broad-band euergv spectrun of accelerated electrous aud protons than voung SNRs (Yamazalietal. 2006)., Old SNRs have quite different characteristics for the expected broad-band energy spectrum of accelerated electrons and protons than young SNRs \citep{yam06}.
+. As an SNR ages. the acceleration time increases. allowing higher maxim euergy of accelerated particles.," As an SNR ages, the acceleration time increases, allowing higher maximum energy of accelerated particles."
+ Towever the losses also increase with time., However the losses also increase with time.
+ The net result is that the electron spectrum is loss-limited when the SNR is older than ~LO? wr. but the proton spectrum ouly jocos loss-limited when the SNR is older than 10? vr.," The net result is that the electron spectrum is loss-limited when the SNR is older than $\sim10^3$ yr, but the proton spectrum only becomes loss-limited when the SNR is older than $\sim10^5$ yr."
+" Therefore. ivan old SNR. the +- aud ταν emissions are expected to be dominated respectively by decay of jeutral pious aud by svuchrotron radiation of secondary electrons from charged. pion decay. respectively,"," Therefore, in an old SNR, the $\gamma$ - and X-ray emissions are expected to be dominated respectively by decay of neutral pions and by synchrotron radiation of secondary electrons from charged pion decay, respectively."
+ These Mons all result from collisions of primary protons ac he ISAL, These pions all result from collisions of primary protons and the ISM.
+ One diagnostic is the ratio of TeV to Nav svuchrotron fluxes., One diagnostic is the ratio of TeV to X-ray synchrotron fluxes.
+ If the protons are accelerated at he shock ruuuiug iuto a CAIC. for example. the ratio should then be consistent with the value obtainec or HESS Jis31-08TS10. (Tianetal.2007).," If the protons are accelerated at the shock running into a GMC, for example, the ratio should then be $\lesssim$ 10, consistent with the value obtained for HESS J1834-087 \citep{tia07}."
+. In. contrast. if the cloud is only illuminated by the particles. the predicted ratio should then be ereater than 100. due to the lack of magnetic field eubliauceimoeut which would occur in the shock case (lence ereatly eubhanciug the ray svuchrotrou emission).," In contrast, if the cloud is only illuminated by the particles, the predicted ratio should then be greater than 100, due to the lack of magnetic field enhancement which would occur in the shock case (hence greatly enhancing the X-ray synchrotron emission)."
+" Based ou the best-fit models (Table 1). the 2-10 keV unabsorbed flux from the rime (iucludiug the three knots) is found to be z 6.9 «10. 1? eres 2 ο,"," Based on the best-fit models (Table 1), the 2-10 keV unabsorbed flux from the ring (including the three knots) is found to be $\approx$ 6.9 $\times$ $^{-12}$ ergs $^{-2}$ $^{-1}$."
+ The sav Buxds Feyπω m Ls teres σα7s + in the TeV band. so the ratio R= E. Εν is ~2.5.," The $\gamma$ -ray flux is $_{\gamma(1-10 {\rm TeV})}$ $\approx$ $\times$ $^{-11}$ ergs $^{-2}$ $^{-1}$ in the 1-10 TeV band, so the ratio $R$ = $_{\gamma}$ $_{X}$ is $\sim$ 2.5."
+ A slightly ower value of δὲ is expected if the X-ray emission from he heavily absorbed westeru-shell is also considered., A slightly lower value of $R$ is expected if the X-ray emission from the heavily absorbed western-shell is also considered.
+ For an old SNR (~ 10! yes for SNR 6353.6-0.7). a heoretical model shows that the fiux iu the 1-10 TeV xuid is a few times higher than that iu 2-10 keV. baud if the TeV enmüssiou is frou piou-decays. the N-ravs are roni svuchrotron emissious of secondary electrou (i the case of a shocked CALC).," For an old SNR $\sim$ $^{4}$ yrs for SNR G353.6-0.7), a theoretical model shows that the flux in the 1-10 TeV band is a few times higher than that in 2-10 keV band if the TeV emission is from pion-decays, the X-rays are from synchrotron emissions of secondary electron (in the case of a shocked GMC)."
+ Svuchrotron enmüssious from secoudazy electrons are about 3 orders of magnitude ueher than that from primary clectrous., Synchrotron emissions from secondary electrons are about 3 orders of magnitude higher than that from primary electrons.
+ Thermal X-ray enission is not considered im this model., Thermal X-ray emission is not considered in this model.
+ Towever. there is little thermal emission originating from au old SNR.," However, there is little thermal emission originating from an old SNR."
+ The new Fermi huge Area Telescope survey released a strong 2-rav source table in the range of 100 MeV - 100 GeV (Abdoetal.2009)., The new Fermi large Area Telescope survey released a strong $\gamma$ -ray source table in the range of 100 MeV - 100 GeV \citep{abd09}.
+. There is no strong GeV source within ΕΠΤσο J17231-317., There is no strong GeV source within HESS J1731-347.
+ For a case that TeV +- rav and X-ray enassions originate from a shocked GAIC. the GeV emission flux is about 2 orders of magnitude lower than that the TeV flux (Fig.," For a case that TeV $\gamma$ -ray and X-ray emissions originate from a shocked GMC, the GeV emission flux is about 2 orders of magnitude lower than that the TeV flux (Fig."
+ 5 of Yamazaki et al.), 5 of Yamazaki et al.)
+ which is much low comparing with other cases (Figs., which is much low comparing with other cases (Figs.
+ 2-G of Yamazaki et al).," 2-6 of Yamazaki et al.),"
+ consistent with the explanation above., consistent with the explanation above.
+ New CO observations support such an SNR scenario: an extended CO cloud at ~ 20 lan | is associated with the SNR., New CO observations support such an SNR scenario: an extended CO cloud at $\sim$ 20 km $^{-1}$ is associated with the SNR.
+" From the spectrum of Sl. we estimate an approximate Io cobluun densitv of ~ [41075 cm7 within the cloud. taking X factor of ~ 2.5 <10°8 Gm? ? ts (Solomon&Barrett1901),"," From the spectrum of S1, we estimate an approximate $_{2}$ column density of $\sim$ $\times$ $^{21}$ $^{-2}$ within the cloud, taking X factor of $\sim$ 2.5 $\times$ $^{20}$ $^{-2}$ $^{-1}$ $^{-1}$ s \citep{sol91}."
+" Asstuuineg the uolecular cloud completely covers the SNR (at least over he extended TESS source region). then it has a mass of My, = Ny,ΟΙ ny 1.) ος ΤΟAZ..."," Assuming the molecular cloud completely covers the SNR (at least over the extended HESS source region), then it has a mass of $M_{H_{2}}$ = $N_{H_{2}}$$\Omega$$d^{2}$ $_{H}$ $M_{\odot})$ $\approx$ $\times$$10^{4}$$M_{\odot}$."
+ This is a GAIC., This is a GMC.
+ Si is located at the peak area of TeV source. aud as a higher local IT). column density than 82 and $3.," S1 is located at the peak area of TeV source, and has a higher local $_{2}$ column density than S2 and S3."
+ This is consistent with the model of a shocked GAIC., This is consistent with the model of a shocked GMC.
+ The evidence of increasing column density from lower-atitude region to higher-latitude reeiou of the SNR revealed bv the CO spectra give a plausible reason ou he XN-ravs distribution., The evidence of increasing column density from lower-latitude region to higher-latitude region of the SNR revealed by the CO spectra give a plausible reason on the X-rays distribution.
+ The ROSAT X-ray nuage of SNR €3523.6-0.7 in the soft baud (0.1-2. keV) appears oulv within the lower-latitude half of the 1radio remmaut where the Ng; is less than 1072 cin7 (e.g. S3 and S2) )ecause absorption blocks soft X-rays but uot hard X- from the upper half where the Ny is lavecr than 1077 ?., The ROSAT X-ray image of SNR G353.6-0.7 in the soft band (0.1-2.4 keV) appears only within the lower-latitude half of the radio remnant where the $_{H}$ is less than $^{22}$ $^{-2}$ (e.g. S3 and S2) because absorption blocks soft X-rays but not hard X-rays from the upper half where the $_{H}$ is larger than $^{22}$ $^{-2}$.
+ From the viewiue-augle of observers. this eives a plysical iuage that the GAIC covers the SNR. so the TeV emissions cover the N-rav and radio emissiou reeions.," From the viewing-angle of observers, this gives a physical image that the GMC covers the SNR, so the TeV emissions cover the X-ray and radio emission regions."
+ The coincidence of the N-xav. morphology with both the radio and TeV 5-ray morphologies sugeests that they are plivsicallvy associated., The coincidence of the X-ray morphology with both the radio and TeV $\gamma$ -ray morphologies suggests that they are physically associated.
+that the average color of Z-EROs ts similar to those of EROS atz-1.5.,that the average color of Z-EROs is similar to those of EROS at $z=1.5$.
+ Therefore. one must interpret Fig.," Therefore, one must interpret Fig."
+ 3 as a plot of the intrinsic color versus (vs.) the redshift of the objects selected as (Z)-EROs/OGs. not as a general characterization of early-type galaxies.," \ref{Fig:colors} as a plot of the intrinsic color versus (vs.) the redshift of the objects selected as (Z)-EROs/OGs, not as a general characterization of early-type galaxies."
+" If we perform a bi-linear fit of the colors as a function of two independent variables ωρι and My. We get with a,=0.69+0.06. αν=-0.02+0.14. a3=0.14+0.06 for the 20 galaxies in this paper."," If we perform a bi-linear fit of the colors as a function of two independent variables $t_{obs.}$ and $M_{V,rest}$ , we get with $a_1=0.69\pm 0.06$, $a_2=-0.02\pm 0.14$, $a_3=0.14\pm 0.06$ for the 20 galaxies in this paper."
+ Given the large error bars. we are unable to discern any gradient of colors with respect to age and metallicity.," Given the large error bars, we are unable to discern any gradient of colors with respect to age and metallicity."
+ L10 did measure significantly non-zero values of a» and «s at z<2.5., L10 did measure significantly non-zero values of $a_2$ and $a_3$ at $z<2.5$.
+ To estimate the average age and stellar mass of our galaxies. we apply the method deseribed in LIO ($4).," To estimate the average age and stellar mass of our galaxies, we apply the method described in L10 4)."
+ We assume a single stellar population for each galaxy: when there is instead a mixture of several populations. our ages would reflect an average of the age for the different populations.," We assume a single stellar population for each galaxy; when there is instead a mixture of several populations, our ages would reflect an average of the age for the different populations."
+ This method uses the synthesis model of Vazdekis et al. (, This method uses the synthesis model of Vazdekis et al. (
+1996). and breaks the age-metallicity degeneracy using the mass—metallicity correlation.,"1996), and breaks the age–metallicity degeneracy using the mass--metallicity correlation."
+ We note that the role of thermally pulsing asymptotic giant branch (TP-AGB) stars is included in this synthesis model. and although some uncertainties are caused by this population (Maraston 2005). these are negligible for 4«6000 (Bruzual 2007): these uncertainties are. however relevant for redder colors.," We note that the role of thermally pulsing asymptotic giant branch (TP-AGB) stars is included in this synthesis model, and although some uncertainties are caused by this population (Maraston 2005), these are negligible for $\lambda < 6000 $ (Bruzual 2007); these uncertainties are, however relevant for redder colors."
+ Moreover. the effect of TP-AGB stars 1s important for ages younger than 1.5 Gyr and. as we later show. our galaxies are probably at this limit.," Moreover, the effect of TP-AGB stars is important for ages younger than 1.5 Gyr and, as we later show, our galaxies are probably at this limit."
+ For galaxies younger than 3 Gyr. the errors in metallicity are insignificant (L10). hence any evolution in the mass-metallicity relationship from high to low redshift would not significantly affect our results.," For galaxies younger than 3 Gyr, the errors in metallicity are insignificant (L10), hence any evolution in the mass-metallicity relationship from high to low redshift would not significantly affect our results."
+ Although the method cannot be applied to each single galaxy separately. it can be used statistically. by binning the galaxies within a small range of redshift and calculating the average age and its r.m.s.," Although the method cannot be applied to each single galaxy separately, it can be used statistically, by binning the galaxies within a small range of redshift and calculating the average age and its r.m.s."
+ for the Z-EROs/OGs in each bin., for the Z-EROs/OGs in each bin.
+ In our case. we divided our sample into three bins. each of which has an associated average intrinsic color (see Fig.," In our case, we divided our sample into three bins, each of which has an associated average intrinsic color (see Fig."
+ 3 and Table 2))., \ref{Fig:colors} and Table \ref{Tab:Candidates_10bins}) ).
+ The age of the early-type galaxies ts plotted in Fig. 4..," The age of the early-type galaxies is plotted in Fig. \ref{Fig:ages},"
+ together with the ages of the galaxies with z<2.5 from L10., together with the ages of the galaxies with $z<2.5$ from L10.
+ The bin with 0.8«z0.9 is not plotted in Fig., The bin with $0.8<z<0.9$ is not plotted in Fig.
+ 4 because it has a very large error bar (L10) and it does not provide any significant information., \ref{Fig:ages} because it has a very large error bar (L10) and it does not provide any significant information.
+ The vertical bars include the errors derived from the uncertainties in the intrinsic colors plus the uncertainties in the synthesis model used in L10., The vertical bars include the errors derived from the uncertainties in the intrinsic colors plus the uncertainties in the synthesis model used in L10.
+ We are neglect the systematic errors in the photometric redshifts: were they not negligible. we would have additional systematic errors in the calculated ages.," We are neglect the systematic errors in the photometric redshifts; were they not negligible, we would have additional systematic errors in the calculated ages."
+ Figure 4. represents the average age of the given sample among the (Z-)ERO/OGs. with all the selection effects associated with each redshift.," Figure \ref{Fig:ages} represents the average age of the given sample among the (Z-)ERO/OGs, with all the selection effects associated with each redshift."
+ In Fig. 5..," In Fig. \ref{Fig:masses},"
+ we plot the stellar masses in each bin derived with the same code., we plot the stellar masses in each bin derived with the same code.
+ The errors in the stellar mass depend on both the data and variations in the initial mass function (IMF) in the Vazdekis et al. (, The errors in the stellar mass depend on both the data and variations in the initial mass function (IMF) in the Vazdekis et al. (
+1996) model. hence we can expect uncertainties of a factor two or even larger.,"1996) model, hence we can expect uncertainties of a factor two or even larger."
+ The stellar masses of the galaxies at z>2.5 are higher than 5x1010 M..., The stellar masses of the galaxies at $z>2.5$ are higher than $5\times 10^{10}$ $M_\odot $.
+ We note that at very high redshift. the stellar masses are much increased. as expected because of a selection effect that prevents us from detecting low mass galaxies at those distances.," We note that at very high redshift, the stellar masses are much increased, as expected because of a selection effect that prevents us from detecting low mass galaxies at those distances."
+ To differentiate between the effects of the dependence on mass and evolution. we apply a bi-linear fit to the average color. weighted with the square inverse of the relative errors in the average age. includingthe z<2.5 bins of LIO.," To differentiate between the effects of the dependence on mass and evolution, we apply a bi–linear fit to the average color, weighted with the square inverse of the relative errors in the average age, includingthe $z<2.5$ bins of L10."
+" We get where both 74; and £s, are in units of Gyr. M. is in units of 10!!Ma. and 6;=0.63+0.04. 5»=0.096+0.008. 0.2]+0.06. c,=-0.18+0.96. c;=1.90+0.29, and cy=4.76+1.77."," We get where both $t_{\rm gal.}$ and $t_{\rm obs.}$ are in units of Gyr, $M_*$ is in units of $10^{11}\ {\rm M_\odot}$, and $b_1=0.63\pm 0.04$, $b_2=0.096\pm 0.008$, $b_3=0.21\pm 0.06$ , $c_1=-0.18\pm 0.96$, $c_2=1.90\pm 0.29$, and $c_3=4.76\pm 1.77$ ."
+" The average epoch of star formation (the first stars might be form earlier) 18 fj,=fobs.—fap. as Shown in Fig. 6.."," The average epoch of star formation (the first stars might be form earlier) is $t_{\rm form.}=t_{\rm obs.}-t_{\rm gal.}$, as shown in Fig. \ref{Fig:ages_form}."
+" Separating the evolution from the mass dependence. where d,=2.18+0.96. d»=-0.9]+0.29. and ds=—4.76+ 1.77."," Separating the evolution from the mass dependence, where $d_1=2.18\pm 0.96$, $d_2=-0.91\pm 0.29$, and $d_3=-4.76\pm 1.77$ ."
+ Forgalaxies 0.8<z2.5 LIO had found that dj=19440.51. d»—0.46+0.32. and d=-0.81 0.98. hence with the extension of the redshift to 3.8 in this paper. we have got with respect to LIO measure values of d and d; that differ significantly from zero.," Forgalaxies $0.8<z<2.5$ L10 had found that $d_1 = 1.94 \pm 0.51$, $d_2 = -0.46 \pm 0.32$, and $d_3 = -0.81 \pm 0.98$ , hence with the extension of the redshift to 3.8 in this paper, we have got with respect to L10 measure values of $d_2$ and $d_3$ that differ significantly from zero."
+ In L10. the stellar masses at 0.8«z<2.5 were confined to à narrow range. although for the wider redshift range 0.8<z«3.8 the range of masses is similarly much larger (see Fig. 5)).," In L10, the stellar masses at $0.8<z<2.5$ were confined to a narrow range, although for the wider redshift range $0.8<z<3.8$ the range of masses is similarly much larger (see Fig. \ref{Fig:masses}) ),"
+ owing to Malmquist bias or other selection effects. permitting a more reliable determination of ds.," owing to Malmquist bias or other selection effects, permitting a more reliable determination of $d_3$."
+" The meaning of our numbers d» and d is that. the lower the redshift (for a fixed mass). the lower the age of galaxy formation. and the higher the stellar mass (for a fixed redshift) the lower the age of galaxy formation,"," The meaning of our numbers $d_2$ and $d_3$ is that, the lower the redshift (for a fixed mass), the lower the age of galaxy formation, and the higher the stellar mass (for a fixed redshift) the lower the age of galaxy formation."
+" On average. for the whole range 0.8<z«3.8. we find that (Sees=—0,91020). (pean=4.76+1.77) Gyr. which is significant at the 3.]—o and 2.7-o levels respectively."," On average, for the whole range $0.8<z<3.8$, we find that $\left \langle \frac{\partial t_{\rm formation}}{\partial t_{\rm observed}}= -0.91\pm 0.29 \right \rangle$, $\left \langle \frac{\partial t_{\rm formation}}{\partial \log _{10}M_*}=-4.76\pm 1.77 \right \rangle $ Gyr, which is significant at the $-\sigma $ and $-\sigma $ levels respectively."
+ We have established a new color criteria for identifying objects of the same intrinsic colors (or even redder) than the EROs/OGs at redshift 0.8<z2.5 but at higher redshift (Z-EROs/OGs)., We have established a new color criteria for identifying objects of the same intrinsic colors (or even redder) than the EROs/OGs at redshift $0.8<z<2.5$ but at higher redshift (Z-EROs/OGs).
+ This is useful for searching for the oldest galaxies at an epoch during which the Universe was younger than ~3 Gyr., This is useful for searching for the oldest galaxies at an epoch during which the Universe was younger than $\sim 3$ Gyr.
+ By adding 20 galaxies at redshifts 2.5<z3.8 to the catalog of LIO. we have been able to see that we still derive formation ages that are within a narrow range (see Fig. 6).," By adding 20 galaxies at redshifts $2.5<z<3.8$ to the catalog of L10, we have been able to see that we still derive formation ages that are within a narrow range (see Fig. \ref{Fig:ages_form}) ),"
+" of on average 2.0+0.3 Gyr (zia,=3— 4) for the galaxies observedat 0.8«z2.5 (LIO). and 0.9705 Gyr [simple weighted average of the three bins in Fig. 6]]."," of on average $2.0\pm 0.3$ Gyr $z_{\rm form.}\approx 3-4$ ) for the galaxies observedat $0.8<z<2.5$ (L10), and $0.9^{+0.4}_{-0.8}$ Gyr [simple weighted average of the three bins in Fig. \ref{Fig:ages_form}] ],"
+" and (ziii,2 4-7. lo)for the galaxies observed at 2.5<z 3.8."," and $z_{\rm form.} \gtrsim 4.7$ , $\sigma$ )for the galaxies observed at $2.5<z<3.8$ ."
+ Clearly. the new galaxies observed at higher redshift were formed (around | Gyr) earlier. mainly because of the higher stellar mass.," Clearly, the new galaxies observed at higher redshift were formed (around 1 Gyr) earlier, mainly because of the higher stellar mass."
+ This is the most astonishing result derived in this paper: that very massive galaxies were formed at redshifts > 4.7., This is the most astonishing result derived in this paper: that very massive galaxies were formed at redshifts $\gtrsim 4.7$ .
+elliplical galaxies. ancl the similarly massive bulges of disc galaxies. regularly contain both a NC and massive black hole (e.g. Graham Driver 2007: GGonzállez Delgado οἱ 22003: seth et 22008).,"elliptical galaxies, and the similarly massive bulges of disc galaxies, regularly contain both a NC and massive black hole (e.g. Graham Driver 2007; Gonzállez Delgado et 2008; Seth et 2008)."
+ Graham Spitler (2009) have quantified how the μι(Vive mass ratio (Fou) increases wilh increasing host spheroid stellar mass. Mj. until only a MBIT is present al (he centre.," Graham Spitler (2009) have quantified how the $M_{\rm BH}/M_{\rm NC}$ mass ratio $F_{\rm BH}$ ) increases with increasing host spheroid stellar mass, $M_{\rm sph}$, until only a MBH is present at the centre."
+ While the runaway merger of NC stars during a merger event may. lead. to (heir conversion into a MBII (e.g.. Zeldovich Podurets 1965: Frank Rees 1976: Quinlan Shapiro 1987; Lee 1993). and feedback processes may also impact £pi (McLaughlin et 22006: Navakshin et 22009). this Letter explores whether dense NCs wil seed MDIIS might evaporate during a collision due to dvnamical heating by the MBIIs.," While the runaway merger of NC stars during a merger event may lead to their conversion into a MBH (e.g., Zel'dovich Podurets 1965; Frank Rees 1976; Quinlan Shapiro 1987; Lee 1993), and feedback processes may also impact $F_{\rm BH}$ (McLaughlin et 2006; Nayakshin et 2009), this Letter explores whether dense NCs with seed MBHs might evaporate during a collision due to dynamical heating by the MBHs."
+ Working from an established N-body code (Bekki et 22004: Dekki 2010) whieh runs on the (Ανν PipE (Sugimoto et 11990). we have developed an idealized model in which a new single NC can be formed [rom the collisionless merger of two NCs with MDIIs an event likely (o occur during a major galaxy merger (e.g.. Dekki 2007a: Bekki οἱ 22010 in preparation).," Working from an established $N$ -body code (Bekki et 2004; Bekki 2010) which runs on the GRAvity PipE (Sugimoto et 1990), we have developed an idealized model in which a new single NC can be formed from the collisionless merger of two NCs with MBHs — an event likely to occur during a major galaxy merger (e.g., Bekki 2007a; Bekki et 2010 in preparation)."
+ Here we investigate how the final structure of the new NC depends on the mass ratio Mpyu/AMxe (= Fygj) of the initial NCs., Here we investigate how the final structure of the new NC depends on the mass ratio $M_{\rm BH}/M_{\rm NC}$ $=F_{\rm BH}$ ) of the initial NCs.
+ We assume that the dynamical evolution of the two NCs are dominated by the NCs and MDIIs themselves. rather than by the gravitational field of backeround stars.," We assume that the dynamical evolution of the two NCs are dominated by the NCs and MBHs themselves, rather than by the gravitational field of background stars."
+ Thus. each of the present models includes only two NCS and two AIBIIs: it ineludes neither background field stars nor external tidal fields ol galaxy mergers.," Thus, each of the present models includes only two NCs and two MBHs: it includes neither background field stars nor external tidal fields of galaxy mergers."
+ The total mass and size of an initial NC are represented by Mc and [νο respectively.," The total mass and size of an initial NC are represented by $M_{\rm NC}$ and $R_{\rm NC}$, respectively."
+ All masses and leneths are measured in units of Ac and Hxc unless otherwise specilied., All masses and lengths are measured in units of $M_{\rm NC}$ and $R_{\rm NC}$ unless otherwise specified.
+ Velocity and time are measured in units of e = (GAMxce/Bxe)? and lay = (Io-”--. respectively. and (he gravitational constant. G is assumed to be 1.," Velocity and time are measured in units of $v$ = $
+(GM_{\rm NC}/R_{\rm NC})^{1/2}$ and $t_{\rm dyn}$ = $(R_{\rm
+NC}^{3}/GM_{\rm NC})^{1/2}$, respectively, and the gravitational constant $G$ is assumed to be 1."
+ If we adopt yc = 5.1 x 105 M. and Rye = 77 pe as fiducial values corresponding tow Cen (e.g.. Mevlan et 11995). which is considered to originate from a nucleated galaxy. (e.g.. Dekki Freeman 2003) (hen e = 16.9 km από {ανν = 446 x 105 vr.," If we adopt $M_{\rm NC}$ = 5.1 $\times$ $10^{6}$ $ \rm M_{\odot}$ and $R_{\rm NC}$ = 77 pc as fiducial values — corresponding to $\omega$ Cen (e.g., Meylan et 1995), which is considered to originate from a nucleated galaxy (e.g., Bekki Freeman 2003) — then $v$ = 16.9 km $^{-1}$ and $t_{\rm dyn}$ = 4.46 $\times$ $10^{6}$ yr."
+" The gravitational softening length. y€,. is set equal to the mean separation of stellar particles at the hall-mass of the initial NC: eg=0.01Axe (=0.77 pe)."," The gravitational softening length, ${\epsilon}_{\rm g}$, is set equal to the mean separation of stellar particles at the half-mass of the initial NC: ${\epsilon}_{\rm g}=0.01R_{\rm NC}$ (=0.77 pc)."
+ This softening length is also adopted for the MDIIs., This softening length is also adopted for the MBHs.
+ The radial densitv profile of our preliminary NC is given by a Plummer model with, The radial density profile of our preliminary NC is given by a Plummer model with
+Although PCA is a mathematically and umuerically robust techuique for analyzing patterus in data. interpreting its results cau be ambiguous.,"Although PCA is a mathematically and numerically robust technique for analyzing patterns in data, interpreting its results can be ambiguous."
+ Iu particular. we would like to verify to what extent there is a one-to-one correspondence between the cigeucolors output ly the PCA aud real surfaces on Earth.," In particular, we would like to verify to what extent there is a one-to-one correspondence between the eigencolors output by the PCA and real surfaces on Earth."
+ To this eud. we test the PCA routine on a suite of simulated data produced v the Virtual Planetary Laboratorys validated 3D spectral Earth model (details cau be found iu Robinson et al.," To this end, we test the PCA routine on a suite of simulated data produced by the Virtual Planetary Laboratory's validated 3D spectral Earth model (details can be found in Robinson et al."
+ 2011)., 2011).
+ The simulations used here were designed to closely πιο the Earthl EPOXI observations taken iu March 2008., The simulations used here were designed to closely mimic the Earth1 EPOXI observations taken in March 2008.
+ We run five different versions of the VPL 3D Earth nodel: 1) Standard: this model is au excellent fit to he EPOXI Earthl observations: the remaining models are identical. but iun cach case a suele model clement das been “turned off: 2) Cloud Free: 3) No Ravleigh Scattering: 1) Black Oceans: 5) Black Laud.," We run five different versions of the VPL 3D Earth model: 1) Standard: this model is an excellent fit to the EPOXI Earth1 observations; the remaining models are identical, but in each case a single model element has been “turned off”: 2) Cloud Free; 3) No Rayleigh Scattering; 4) Black Oceans; 5) Black Land."
+ We show he results of this experiment in Appendix I: here we sinplv state our conclusions: 1) PCA successfully determines the dimoeusionalitv of he color variability aud therefore the ΠΠ ΓΕuimniber of different surface types coutributing to color variations., We show the results of this experiment in Appendix I; here we simply state our conclusions: 1) PCA successfully determines the dimensionality of the color variability and therefore the minimum number of different surface types contributing to color variations.
+ Iu particular. -dimensional variations require 11l surface types (N.B. we count clouds as a surface type).," In particular, $n$ -dimensional variations require $n+1$ surface types (N.B. we count clouds as a surface type)."
+ 2) Ravleigh scattering is iuportaut in cleterumuine the time-averaged broadbaud colors of Earth. but does not sienificautly affect its rotational color variability.," 2) Rayleigh scattering is important in determining the time-averaged broadband colors of Earth, but does not significantly affect its rotational color variability."
+ 3) Cloud-free land. surfaces. which are red. contribute a red eigeucolor to the diurnal variability.," 3) Cloud-free land surfaces, which are red, contribute a red eigencolor to the diurnal variability."
+ The presence of relatively cloud-free land (deserts) near the equator explains why the rotational map of the red cigencolor (Figure10inCowanetal.2009). successtully identified the major landforms aud bodies of water ou Earth., The presence of relatively cloud-free land (deserts) near the equator explains why the rotational map of the red eigencolor \citep[Figure~10 in][]{Cowan_2009} successfully identified the major landforms and bodies of water on Earth.
+ 1) Oceanus ire essentially a null surface. coutributing ucither to the broadband colors of Earth. nor to the time-variabilitv of those colors. except insofar as the presence of oceans corresponds to a shortage of laud.," 4) Oceans are essentially a null surface, contributing neither to the broadband colors of Earth, nor to the time-variability of those colors, except insofar as the presence of oceans corresponds to a shortage of land."
+" 5) In the absence of laud. the variability is οταν, due to large-scale inhomogeneities m cloud cover."," 5) In the absence of land, the variability is gray, due to large-scale inhomogeneities in cloud cover."
+ 6) PCA necessarily outputs orthogonal cigencolors aud a good deal of Earth's variability is due to clouds., 6) PCA necessarily outputs orthogonal eigencolors and a good deal of Earth's variability is due to clouds.
+ Therefore. if the first eigencolor is red. then the second eieencolor nav be blue even if there is no blue surface rotating iu and out of view: this is an muprovenieut on the interpretation of Cowanetal.(2009).," Therefore, if the first eigencolor is red, then the second eigencolor may be blue even if there is no blue surface rotating in and out of view; this is an improvement on the interpretation of \cite{Cowan_2009}."
+. Iu Figures d. 5 we show the eigeuvalue spectra for time-variations in the 7 cigencolors identified by the PCA of the EPOXI polar observations., In Figures \ref{polar1_variability} \ref{polar2_variability} we show the eigenvalue spectra for time-variations in the 7 eigencolors identified by the PCA of the EPOXI polar observations.
+ The cigenvalue for a even component is the projection of the data variance outo that cigeuvector: we plot here the square roots of the eieenvalues. which is a measure of the RAIS variability of the data projected onto an eigeuvector.," The eigenvalue for a given component is the projection of the data's variance onto that eigenvector; we plot here the square root of the eigenvalues, which is a measure of the RMS variability of the data projected onto an eigenvector."
+ The variability has been normalized in the figures such that the suu of the variability for all seven comupoucuts is unity., The variability has been normalized in the figures such that the sum of the variability for all seven components is unity.
+ By definition. the low-order principal components have the largest variance.," By definition, the low-order principal components have the largest variance."
+ For the North observation. there are two cigcucolors that dominate the color variatious of Earth: the third eieencolor contributes only —L% of the planets color variability.," For the North observation, there are two eigencolors that dominate the color variations of Earth: the third eigencolor contributes only $\sim4$ of the planet's color variability."
+ As in Cowanetal.(2009)... this means that the colors of Earth populate a two-dimensional plane rather than fillme the eutire seveu-dineusional color-space. and this requires at least three surface types.," As in \cite{Cowan_2009}, this means that the colors of Earth populate a two-dimensional plane rather than filling the entire seven-dimensional color-space, and this requires at least three surface types."
+ The southern observation. on the other haud. is dominated by a sinele eieencolor (the second cigencolor coutributes to variability at the <10% level).," The southern observation, on the other hand, is dominated by a single eigencolor (the second eigencolor contributes to variability at the $<10$ level)."
+ This meaus that for the 21 hours of observations— the colors of the planet populated a one-dimensional line iu the seven-dimensional color vole. requiring ouly two surface types.," This means that —for the 24 hours of observations— the colors of the planet populated a one-dimensional line in the seven-dimensional color volume, requiring only two surface types."
+ The cigencolors (the AA) from Equation 1)) are shown in Figures 6 7.., The eigencolors (the $A_{i}(\lambda)$ from Equation \ref{pca}) ) are shown in Figures \ref{polar1_eigenspectra} \ref{polar2_eigenspectra}.
+ The raw cigeucolors are by definition— orthogonal and normalized ($75 1). aud this is how we preseuted them in Cowanetal. (2009)..," The raw eigencolors are —by definition— orthogonal and normalized $\sum_{j=1}^{7}A_{i}^{2}(\lambda_{j})=1$ ), and this is how we presented them in \cite{Cowan_2009}. ."
+ Tere we have instead scaled the cigencolors by, Here we have instead scaled the eigencolors by
+system obtained by Adéuetal.(2009).,system obtained by \citet{Aden2009}.
+. We fiud at the mass within the volume euclosed by our ≻↴⋝↴∖↴↸∖↥⋅↖↽↸∖≺↧↴∖↴↑⋜∐⋅↴∖↴↕↴∖↴⋅≩∙∣↴⊥⊓∖↓∣↙⇂⋀∐∙∙↕↸∖⋜∥∐∐∶↴∙↑∪⋜↧ πο” ⋅⋡ uass-to-lielit ratio of103 ο[ALΤΙ. Interestingly. ∐∖↕⊔⋜↧↴∖↴↴∖↴↖↖↽↕↑∐↕∐∶≩∩," We find that the mass within the volume enclosed by our observed stars is $3.7^{+2.2}_{-1.6}\times \, 10^6 \, M_{\odot}$, leading to a mass-to-light ratio of $103^{+83}_{-48} \, [M_{\odot}/L_{\odot}]$."
+∩∏≻↸⊳↕↴∖↴↴∖↴↕∶↴∙⊾∐∐↸⊳⋜↧∐∏⋅↖↽↕∪↖↖⇁↸∖↥⋅↑∐⋜⋯ ∐∖↘⇁↸⊳∪∐∐⊔∪∐⋯⋜↧↴∖↴↴∖↴↴∖↴↸⊳⋜↧↕↸∖⋅⊲↕≯≺∏⋯≺∏⋝∙↖↽≋⊓⋅↕∶↴∙⊾⋜∐⋅↕↸∖↑⋜↕↕∙ ⊇∩∩≺∖∖⋟∙∙↴∖↴∏∶↴∙⊾∶↴∙⊾↸∖↴∖↴↑↕∐∶↴⋁∐⋜↧↑∐↸∖↥⋅↸⊳∏↕↸∖↴∖↴≼↧∪↸∖↴∖↴∐∪↑↴∖↴∐⋜⋯∖ 16 halo properties seen in other dSplis.," Interestingly, the mass within pc is significantly lower than the “common mass scale” found by \citet{2008Natur.454.1096S}, suggesting that Hercules does not share the halo properties seen in other dSphs."
+ We found teutative evidence for a velocity eradient of 163klauss !kpe t. and evidence of an asviunietrie exteusion in the light distribution at 0.5 kkpc.," We found tentative evidence for a velocity gradient of $16 \pm
+3$ $^{-1}$ $^{-1}$, and evidence of an asymmetric extension in the light distribution at $\sim0.5$ kpc."
+ We explored the hypothesis tha these features are du| fo tidal iuteractions with the Milkv. Wary., We explored the hypothesis that these features are due to tidal interactions with the Milky Way.
+" AÁssunüueg a tidal radius of 185 ppc. we show that a self-consisteut mode requires Hercules to be ou an orbit with periceutre ry=Lot and with a mass within r; of Mair—52DE$5 LOPN,"," Assuming a tidal radius of $485$ pc, we show that a self-consistent model requires Hercules to be on an orbit with pericentre $r_p = 18.5
+\pm 5$ kpc, and with a mass within $r_t$ of $M_{\mathrm{tid},r_t}=5.2_{-2.7}^{+2.7}\times10^6$ $_\odot$."
+ We are very erateful to Matt Walker for usefu discussions regarding the mass estimates im earlier works., We are very grateful to Matt Walker for useful discussions regarding the mass estimates in earlier works.
+ We thank Pavel Ivoupa for stimulating discussions., We thank Pavel Kroupa for stimulating discussions.
+ D.À thanks Lennart Lindeeren at Lund Observatory for help with statistics., D.A thanks Lennart Lindegren at Lund Observatory for help with statistics.
+ S.F. is a Roval Swedish: Academy. of Scicuces Research Fellow supported by a eraut from the dEAuut and Alice Wallenbere Foundation., S.F. is a Royal Swedish Academy of Sciences Research Fellow supported by a grant from the Knut and Alice Wallenberg Foundation.
+ M.LW is supported by a Roval Society University Research Fellowship., M.I.W is supported by a Royal Society University Research Fellowship.
+the kinetic equation contains only the term qualitatively similar to the first term in equation (7). the maximum scattering rate is in the direction antiparallel to the particle velocity. and the photons shift monotonically toward lower frequencies.,"the kinetic equation contains only the term qualitatively similar to the first term in equation (7), the maximum scattering rate is in the direction antiparallel to the particle velocity, and the photons shift monotonically toward lower frequencies."
+ The solution of the svstem (S) is given hy One can see that the elliclency of intensity transfer between the beams is determined. by Lofar., The solution of the system (8) is given by One can see that the efficiency of intensity transfer between the beams is determined by $\Gamma\equiv Iar$.
+ At laree enough DI. {νι:Q and {νι-2.," At large enough $\Gamma$ , $I_{\nu_a}\to 0$ and $I_{\nu_b}\to I$."
+" Since the initial intensity of pulsar beam greatly. exeeeds that of. the background s ↓⋅⋯⊔⋜∐↓∪⊔⊳∫⊧UpTIED∣∖≮∖↕↓⊳⇂↓↕∢⊾↓⊔↿∢⊾⊔⊳∖∐∙∖⇁↿↓⋅⋜⋯⊳∖⇂⋖⋅↓⋅↿∪↿⇂↥⋖⋅ii- . ⋅ background is significant on condition that Note that P includes the total intensity Z. Le. actually Lh,Ohi and lor laree enough E the scattering is ellicient independently of the smallness of the backgroundintensity."," Since the initial intensity of pulsar beam greatly exceeds that of the background radiation, $I_{\nu_b}^{(0)}/I_{\nu_a}^{(0)}\ll 1$, the intensity transfer to the background is significant on condition that Note that $\Gamma$ includes the total intensity $I$, i.e. actually $I_{\nu_a}^{(0)}$, and for large enough $\Gamma$ the scattering is efficient independently of the smallness of the backgroundintensity."
+ Now we are to estimate the scattering ellicicney. EP—Lar. where foἐνIN and @ is given by equation (9). for the parameters relevant to pulsar magnetosphere.," Now we are to estimate the scattering efficiency, $\Gamma =Iar$, where $I\approx I_{\nu_a}^{(0)}$ and $a$ is given by equation (9), for the parameters relevant to pulsar magnetosphere."
+" The spectrum. of pulsar radiation generally has the power-law forni. £4""HGfo) ""C. and the spectral intensity at frequencies Mllz is related to the total raclio Luminosity of a pulsar. L.as l4,=Lftys. where 8=ziufd ds the cross-section of the pulsar. bean at a distance row the pulse width in the angular measure."," The spectrum of pulsar radiation generally has the power-law form, $I_{\nu_a}^{(0)}=I_{\nu_0}(\nu_a/\nu_0)^{-\alpha}$ , and the spectral intensity at frequencies $\nu_0\sim 100$ MHz is related to the total radio luminosity of a pulsar, $L$, as $I_{\nu_0}=L/\nu_0S$, where $S=\pi r^2w^2/4$ is the cross-section of the pulsar beam at a distance $r$, $w$ the pulse width in the angular measure."
+ Ht is convenient to normalize the number density of the scattering particles to the Coldreich-Julian number density. 7.=8DPec. where & is the plasma multiplicity factor. ancl £2? is the pulsar period.," It is convenient to normalize the number density of the scattering particles to the Goldreich-Julian number density, $n_e=\kappa B/Pce$, where $\kappa$ is the plasma multiplicity factor, and $P$ is the pulsar period."
+" Weeping in mind that n,xBoxr.. one can estimate the scattering elliciencev as where D, is the magnetic field strength at the stellar surface. and it is taken that the neutron star radius is 10"" em."," Keeping in mind that $n_e\propto B\propto r^{-3}$, one can estimate the scattering efficiency as where $B_\star$ is the magnetic field strength at the stellar surface, and it is taken that the neutron star radius is $10^6$ cm."
+ ALL the quantities in equation (13) are normalized to their tvpical values., All the quantities in equation (13) are normalized to their typical values.
+ One can see that the scattering ellicicney may be as large. as about a few tens., One can see that the scattering efficiency may be as large as about a few tens.
+ To conclude whether the condition of ellicient intensity transfer given byequation. (12) can indeed be satisfied: let us estimate the level of background. radiation resulting [rom the spontaneous scattering. of. the pulsar radio. beam. iU! πασά.," To conclude whether the condition of efficient intensity transfer given byequation (12) can indeed be satisfied let us estimate the level of background radiation resulting from the spontaneous scattering of the pulsar radio beam, $I_{\nu_b}^{(0)}\sim I_{\nu_a}^{(0)}n_e\eta r{\rm
+d}\sigma/{\rm d}\Omega_1$ ."
+" ""Taking into account that de/dO,=sin?66,(5 gp). ητε67/2 and pnl77. we find that llence. typically IUUnHUSνο1ο710£7. and to salisfv equation (12) P=1528 are necessary."," Taking into account that ${\rm d}\sigma
+/{\rm
+d}\Omega_1=r_e^2\sin^2\theta\sin^2\theta_1/(\gamma^6\eta_1^4\eta^2)$ , $\eta\approx \theta^2/2$ and $\eta_1\approx 1/\gamma^2$, we find that Hence, typically $I_{\nu_b}^{(0)}/I_{\nu_a}^{(0)}\sim
+10^{-8}-10^{-12}$, and to satisfy equation (12) $\Gamma=18-28$ are necessary."
+ Thus. the intensity transfer from the radio beam to the background can indeed be significant in pulsars. especially at low enough frequencies.," Thus, the intensity transfer from the radio beam to the background can indeed be significant in pulsars, especially at low enough frequencies."
+ Let us consider the location of the scattering region in the magnetosphere of a pulsar., Let us consider the location of the scattering region in the magnetosphere of a pulsar.
+ According to equation (13). Po shows strong explicit dependence on r. Dx+D ," According to equation (13), $\Gamma$ shows strong explicit dependence on $r$, $\Gamma\propto r^{-4}$."
+Llowever. the implicit dependence on r appears still more significant.," However, the implicit dependence on $r$ appears still more significant."
+ As 6xr (see below). ο=νι0xp ," As $\theta\propto r$ (see below), $\nu_a=\nu_b/\theta^2\gamma^2\propto r^{-2}$."
+This can be understood. as follows., This can be understood as follows.
+ At a fixed. [frequency μι. the background radiation actually grows because of the scattering. of dillerent. radio beam frequencies £6. which satisly the condition £i=(6(r)s7 at different altitudes r.," At a fixed frequency $\nu_b$, the background radiation actually grows because of the scattering of different radio beam frequencies $\nu_a$, which satisfy the condition $\nu_b=\nu_a\theta^2(r)\gamma^2$ at different altitudes $r$."
+ Larger altitudes imply. lower frequencies £7. in which case the incident intensity of the radio beam is much Larger and stimulates stronger scattering.," Larger altitudes imply lower frequencies $\nu_a$, in which case the incident intensity of the radio beam is much larger and stimulates stronger scattering."
+ Taking into account the above considerations. one can obtain that Dxi77.0.," Taking into account the above considerations, one can obtain that $\Gamma\propto r^{2\alpha
+-4}$."
+ Hence. the scattering elliciency increases with distance at à2 and decreases at à«2.," Hence, the scattering efficiency increases with distance at $\alpha>2$ and decreases at $\alpha<2$."
+ At high enough frequencies. =1 Cle. three of the pulsars with the precursor components. namely the Crab. the Vela and PSR. DIS22-09. have à=2.8. 2.7 ancl 2.3. respectively. whereas for PSR. DBI055-52 there is no spectral data in this region (perhaps. the spectrum is too steep for the pulsar to be detectable).," At high enough frequencies, $\ga 1$ GHz, three of the pulsars with the precursor components, namely the Crab, the Vela and PSR B1822-09, have $\alpha=2.8$, 2.7 and 2.3, respectively, whereas for PSR B1055-52 there is no spectral data in this region (perhaps, the spectrum is too steep for the pulsar to be detectable)."
+ At lower frequencies such a steep spectrum is preserved only in the Crab pulsar. while in the other pulsars under consideration à drops below 2.," At lower frequencies such a steep spectrum is preserved only in the Crab pulsar, while in the other pulsars under consideration $\alpha$ drops below 2."
+ Note. however. that at lower frequencies the scattering is more efficient and may noticeably suppress the radio beam. intensity. leading to the pulsar spectrum fattening. so that the original spectra may be much steeper than the observed ones.," Note, however, that at lower frequencies the scattering is more efficient and may noticeably suppress the radio beam intensity, leading to the pulsar spectrum flattening, so that the original spectra may be much steeper than the observed ones."
+" Thus. we assume that D increases with altitude and. corresponcdinelv. the induced scattering is most ellicient. at the upper boundary of the scattering region. ic. at distances of order of the evelotron resonance radius. which is defined as 235,507/2=we: and estimated as where r; is the σαι evlinder radius and it is taken that @=rí2rg (see below)."," Thus, we assume that $\Gamma$ increases with altitude and, correspondingly, the induced scattering is most efficient at the upper boundary of the scattering region, i.e. at distances of order of the cyclotron resonance radius, which is defined as $2\pi\nu_a\gamma\theta^2/2=\omega_G$ and estimated as where $r_L$ is the light cylinder radius and it is taken that $\theta=r/2r_L$ (see below)."
+ One can see that the region of evcelotron resonance tvpicallv lies in the outer magnetosphere. at distances of order of the light. exlinder radius.," One can see that the region of cyclotron resonance typically lies in the outer magnetosphere, at distances of order of the light cylinder radius."
+ The location of the scattered. component in the pulse profile can be examined. as follows., The location of the scattered component in the pulse profile can be examined as follows.
+ Since the scattering region lies well above the emission region. the wavevector Kk makes the angle ~παςτε with the instantaneous direction of the magnetic axis of the rotating magnetosphere (here ὁ is the angle between the rotational ancl magnetic axes of a pulsar).," Since the scattering region lies well above the emission region, the wavevector ${\bmath k}$ makes the angle $\sim r\sin\zeta /r_L$ with the instantaneous direction of the magnetic axis of the rotating magnetosphere (here $\zeta$ is the angle between the rotational and magnetic axes of a pulsar)."
+ As the polar angle. of the point. of scattering is μαςπε. in the dipolar geometry the local magnetic field vector 6 is inclined at the angle 3rπαςδι to the magnetic axis. so that the angle between & and b is rsinc/2r;.," As the polar angle of the point of scattering is $r\sin\zeta
+/r_L$, in the dipolar geometry the local magnetic field vector ${\bmath b}$ is inclined at the angle $3r\sin\zeta /2r_L$ to the magnetic axis, so that the angle between ${\bmath k}$ and ${\bmath
+b}$ is $r\sin\zeta /2r_L$."
+ In the corotating frame. the scattered radiationis directed. approximately along 5.," In the corotating frame, the scattered radiationis directed approximately along ${\bmath b}$ ."
+" Then in the laboratory frame it is shifted. because of rotational aberration by the angle παςsr, in the direction of rotation. Le. toward the magnetic axis."," Then in the laboratory frame it is shifted because of rotational aberration by the angle $r\sin\zeta /r_L$ in the direction of rotation, i.e. toward the magnetic axis."
+ Thus. the scattered component. precedes the main pulse by AXA~rsinc/2r; in longitude.," Thus, the scattered component precedes the main pulse by $\Delta\lambda\sim r\sin\zeta
+/2r_L$ in longitude."
+ One can see that the precursor. separation from the main pulse is determined. by the height of the scattering region. r. and," One can see that the precursor separation from the main pulse is determined by the height of the scattering region, $r$ , and"
+is a joint project of (he University of Massachusetts aud Infrared. Processing anc Analysis Center/California. Institute of Technology. founded by the National Aeronauties and Space Administration and the National Science Foundation.,"is a joint project of the University of Massachusetts and Infrared Processing and Analysis Center/California Institute of Technology, founded by the National Aeronautics and Space Administration and the National Science Foundation."
+complete.,complete.
+" The fit is done in two ways: fitting all redshift bins simultaneously to estimate the global evolution of the LF at 0.0€z<0.6, and fitting each bin individually to explore any bin-specific deviations from the global law."," The fit is done in two ways: fitting all redshift bins simultaneously to estimate the global evolution of the LF at $0.0 \leq z \leq 0.6$, and fitting each bin individually to explore any bin-specific deviations from the global law."
+" We also integrate the individually fitted LFs over luminosity to determine the luminosity density, which will be used to estimate the star formation rate density Section ??))."," We also integrate the individually fitted LFs over luminosity to determine the luminosity density, which will be used to estimate the star formation rate density Section \ref{sec:IR_LF_SFR}) )."
+" As mentioned in Section ??,, the LLF is affected by the"," As mentioned in Section \ref{sec:IR_LF_Method}, the LLF is affected by the"
+on the line profiles we computed.,on the line profiles we computed.
+ In Sect.4.. we show how MCRT is able to reproduce typical BAL QSOs line profiles.," In \ref{fitbal}, we show how MCRT is able to reproduce typical BAL QSOs line profiles."
+ We discuss the results of the line profile fitting and summarize our conclusions in the last two sections of the paper., We discuss the results of the line profile fitting and summarize our conclusions in the last two sections of the paper.
+ MCRT ts a Fortran77 fully 3D Monte Carlo (MC) radiative transfer (RT) code that we built to compute the resonance line profiles produced in axisymmetric winds., MCRT is a Fortran77 fully 3D Monte Carlo (MC) radiative transfer (RT) code that we built to compute the resonance line profiles produced in axisymmetric winds.
+ The use of the Monte Carlo simulation technique allows the radiative transfer equation to be solved exactly (i.e. without making use of the Sobolev approximation). as well as ensuring the self consistent treatment of the radiative coupling between distant regions in a wind subject to more complex velocity fields than monotonic radial laws (e.g. Knigge et al. 1995)).," The use of the Monte Carlo simulation technique allows the radiative transfer equation to be solved exactly (i.e. without making use of the Sobolev approximation), as well as ensuring the self consistent treatment of the radiative coupling between distant regions in a wind subject to more complex velocity fields than monotonic radial laws (e.g. Knigge et al. \cite{kn95}) )."
+ Monte Carlo RT code have been extensively described (e.g. Knigge et al. 1905..," Monte Carlo RT code have been extensively described (e.g. Knigge et al. \cite{kn95},"
+ Wood et al. 2001.. ," Wood et al. \cite{wo01}, ,"
+Dijkstra et al. 2006)).," Dijkstra et al. \cite{di06}) ),"
+ so that we only recall here the fundamental principles of this technique and the particularities of the code we developed., so that we only recall here the fundamental principles of this technique and the particularities of the code we developed.
+ As stated in the introduction. our main goal is to identify of the key ingredients (geometry and overall kinematics) of the wind governing the typical profile of the BAL QSO UV resonance lines.," As stated in the introduction, our main goal is to identify of the key ingredients (geometry and overall kinematics) of the wind governing the typical profile of the BAL QSO UV resonance lines."
+" Thus we do not consider negligible effects. such as the relativistic ones that remain small even for outflows with high (yay< 0.2c) terminal speed (Hutsemékkers Surdej 1990)) or the fact that the line is a resonance doublet. because the velocity separation of the doublet components is small with respect to vj, (e.g. Hewitt et al. 1974.."," Thus we do not consider negligible effects, such as the relativistic ones that remain small even for outflows with high $\rm{v}_{max} \leq 0.2 c$ ) terminal speed (Hutsemékkers Surdej \cite{hu90}) ) or the fact that the line is a resonance doublet, because the velocity separation of the doublet components is small with respect to $\rm{v}_{max}$ (e.g. Hewitt et al. \cite{he74},"
+ Grinin 1984)., Grinin \cite{gr84}) ).
+ When using the MC technique. the solution of the RT equation is found by following a huge number of photons on their way through the wind.," When using the MC technique, the solution of the RT equation is found by following a huge number of photons on their way through the wind."
+ Each step in the photon’s life (position and direction of emission. position of interaction. ete) 1s determined by the mean of random numbers distributed according to the normalized probability density function (NPDF) of the corresponding simulated physical process.," Each step in the photon's life (position and direction of emission, position of interaction, etc) is determined by the mean of random numbers distributed according to the normalized probability density function (NPDF) of the corresponding simulated physical process."
+ Thus if one wishes that the frequency v; of all of the emitted photons follows a given law L(v) over the frequency interval [νιως]. then v; will be randomly chosen by solving the transformation equation (Press et al. 1992)):," Thus if one wishes that the frequency $\nu_i$ of all of the emitted photons follows a given law $L(\nu)$ over the frequency interval $[\nu_{min},\nu_{max}]$, then $\nu_i$ will be randomly chosen by solving the transformation equation (Press et al. \cite{pr92}) ):"
+ where € is a random number drawn from a uniform distribution in the interval [0.1].," where $\xi$ is a random number drawn from a uniform distribution in the interval $[0,1]$."
+" In MCRT this number is generated using the ""ran2"" subroutine of Press et al. (1992))."," In MCRT this number is generated using the “ran2"" subroutine of Press et al. \cite{pr92}) )."
+ In the following. each new occurrence of & refers to the call of such a new random number.," In the following, each new occurrence of $\xi$ refers to the call of such a new random number."
+" In general. there is seldom an analytical solution to Eq.]. so we implemented the ""table lookup method"" (see Avery House 1968)). which allows us arbitrary NPDF's."," In general, there is seldom an analytical solution to \ref{transfo} so we implemented the “table lookup method"" (see Avery House \cite{av68}) ), which allows us arbitrary NPDF's."
+" In MCRT the initial position of emission of the photons 1s chosen isotropically on the surface of the continuum emission region. which 1s modeled by a sphere of radius Rj, and of infinite optical depth έτο=οὐ) located at the center of the wind."," In MCRT the initial position of emission of the photons is chosen isotropically on the surface of the continuum emission region, which is modeled by a sphere of radius $R_{in}$ and of infinite optical depth $\tau_C = \infty$ ) located at the center of the wind."
+ The direction of travel through the wind is then determined by randomly sampling a half sphere taking into account that the photons are forced to leave the continuum source upward., The direction of travel through the wind is then determined by randomly sampling a half sphere taking into account that the photons are forced to leave the continuum source upward.
+" The wind is filled with 2-level atoms whose rest-frame normalized absorption profile @,,, is described by a Gaussian (Natta Beckwith 1986.. Knigge et al. 1995))"," The wind is filled with 2-level atoms whose rest-frame normalized absorption profile $\phi_{abs}$ is described by a Gaussian (Natta Beckwith \cite{na86}, Knigge et al. \cite{kn95}) )"
+ such that where vo is the rest-frame frequency of the considered transition. and in which A ensures a continuous transition between the absorption profile and the zero intensity.," such that where $\nu_0$ is the rest-frame frequency of the considered transition, and in which $K$ ensures a continuous transition between the absorption profile and the zero intensity."
+" Here. H is a constant allowing for the normalization of the line profile over the interval vo+Αν]. where Av,p, is the full width at zero intensity (FWZLI) of the absorption profile."," Here, $H$ is a constant allowing for the normalization of the line profile over the interval $\nu_0 \pm |\Delta \nu_{abs}|/2$, where $\Delta \nu_{abs}$ is the full width at zero intensity (FWZI) of the absorption profile."
+" We choose the value of Avypy to ensure the continuity of the absorption profile at the border of the interval vo+[|Av;,|/2.", We choose the value of $\Delta \nu_{abs}$ to ensure the continuity of the absorption profile at the border of the interval $\nu_0 \pm |\Delta \nu_{abs}|/2$ .
+" The parameter ο=Av,p/(2N21n2) is such that Avis=2vo(Vn/e) Is the FWHM of the absorption profile.", The parameter $\sigma_{turb}=\Delta \nu_{turb}/(2 \sqrt{2 \ln{2}})$ is such that $\Delta \nu_{turb}=2~\nu_0 (\rm{v}_{turb}/c)$ is the FWHM of the absorption profile.
+ The velocity γην includes the thermal and the macroscopic turbulence components in the wind that broaden the absorption profile., The velocity $\rm{v}_{turb}$ includes the thermal and the macroscopic turbulence components in the wind that broaden the absorption profile.
+ We assume vii to be constant throughout the wind., We assume $\rm{v}_{turb}$ to be constant throughout the wind.
+" Owing to the velocity field V(7.9.9) present in the wind. the initial frequency v; of a photon flying in the direction 7? is seen Doppler-shifted by an atom of the wind in such à way that its “local” frequency v, in the atom rest-frame is given by Thus à photon will enter in resonance with the surrounding atoms only if its local frequency fulfills the condition defining the so-called ""resonance zone: When the photon enters such a region, the opacity of the medium becomes nonzero as does the probability of being absorbed."," Owing to the velocity field $\overrightarrow{\rm{v}}(r,\theta,\phi)$ present in the wind, the initial frequency $\nu_i$ of a photon flying in the direction $\overrightarrow{n}$ is seen Doppler-shifted by an atom of the wind in such a way that its “local"" frequency $\nu_l$ in the atom rest-frame is given by Thus a photon will enter in resonance with the surrounding atoms only if its local frequency fulfills the condition defining the so-called “resonance zone"": When the photon enters such a region, the opacity of the medium becomes nonzero as does the probability of being absorbed."
+" If 7 resonance zones are found along the direction of propagation of the photon in the wind. the total optical depth r,, seen by the photon until it escapes the wind is simply computed as where «4,, 1s the total absorption coefficient. of the considered resonance transition. and a; and 5; are respectively the coordinates of the beginning and of the end of the j resonance zone found along the line of flight of the photon."," If $n$ resonance zones are found along the direction of propagation of the photon in the wind, the total optical depth $\tau_{tot}$ seen by the photon until it escapes the wind is simply computed as where $\kappa_{\nu_0}$ is the total absorption coefficient of the considered resonance transition, and $a_j$ and $b_j$ are respectively the coordinates of the beginning and of the end of the $j^{th}$ resonance zone found along the line of flight of the photon."
+ Given the probabilistic interpretation of the RT. a photon experiencinga total optical depth of τιν) has a probability pESgos of escaping the medium without being absorbed.," Given the probabilistic interpretation of the RT, a photon experiencinga total optical depth of $\tau_{tot}(\nu_i)$ has a probability $p=e^{-\tau_{tot}(\nu_i)}$ of escaping the medium without being absorbed."
+ Thisinterpretation is used in the MC code to identify the, Thisinterpretation is used in the MC code to identify the
+Furthermore. the GeV 5-rays from the SNR. which is known to be interacting with a molecular cloud. with a luminosity of ~11075 erg 3 for a distance of 6 kpe. have been detected. with the LAT on board the Gamma-ray Space Telescope CXbdoctal.2009b).,"Furthermore, the GeV $\gamma$ -rays from the SNR, which is known to be interacting with a molecular cloud, with a luminosity of $\sim1\times10^{36}$ erg $^{-1}$ for a distance of 6 kpc, have been detected with the LAT on board the Gamma-ray Space Telescope \citep[][]{Abea09b}."
+. Ehe 5-rav spectrum cannot be fitted. with a single power law and steepens above a few GeV. Abdoetal.(20095) discussed. the nonthermal radiative properties of the radio emission. and 5-ravs from the SNR using a broken power-law with the same index for the momentum cistribution of the radiating electrons/protons., The $\gamma$ -ray spectrum cannot be fitted with a single power law and steepens above a few GeV. \citet[][]{Abea09b} discussed the nonthermal radiative properties of the radio emission and $\gamma$ -rays from the SNR using a broken power-law with the same index for the momentum distribution of the radiating electrons/protons.
+ Pheir result shows that p-p interaction is more likely the main process to produce the observed 5-rays since the observed radio svnchrotron spectrum cannot be well reproduced in the bremsstrahlung-dominated. case and an origin of inverse Compton scattering for the 5-rays needs a low density of hvdrogen (<0.1 1) in the SNR. which conlliets with the constraints from observations in the radio and N-rav. bancs.," Their result shows that p-p interaction is more likely the main process to produce the observed $\gamma$ -rays since the observed radio synchrotron spectrum cannot be well reproduced in the bremsstrahlung-dominated case and an origin of inverse Compton scattering for the $\gamma$ -rays needs a low density of hydrogen $<0.1$ $^{-1}$ ) in the SNR, which conflicts with the constraints from observations in the radio and X-ray bands."
+ Semi-analvtical methods to the nonlinear dilIusive shock acceleration process have been widely used to investigate the nonthermal radiative. properties of SNRs (e.g..Morlinoctal.2009:FangEllison 2010).," Semi-analytical methods to the nonlinear diffusive shock acceleration process have been widely used to investigate the nonthermal radiative properties of SNRs \citep[e.g.,][]{Mea09,Fea09,Eea10}."
+. In this paper. we study the nonthermal radiative properties of the SNR W51C€ based on a semi-analytical method to the nonlinear diffusive shock acecleration mechanism.," In this paper, we study the nonthermal radiative properties of the SNR W51C based on a semi-analytical method to the nonlinear diffusive shock acceleration mechanism."
+ In our model. a part of the SNI shell transports into the molecular cloud (AIC). then the shock is greatly decelerated due to the large density of matter in the MC.," In our model, a part of the SNR shell transports into the molecular cloud (MC), then the shock is greatly decelerated due to the large density of matter in the MC."
+ As a result. the spectra of the accelerated protons and electrons are determined in a case with a low Mach number.," As a result, the spectra of the accelerated protons and electrons are determined in a case with a low Mach number."
+ We find that the observed. 5-ray. spectrum for W51C€ can be well reproduced: via p-p collisions of the accelerated: protons in the part of shell interacting with the MC. ancl the raclio emission from the SNR can be explained as the radiation of the electrons accelerated by the other part of the SNIt shock freely expanding into the ambient interstellar space.," We find that the observed $\gamma$ -ray spectrum for W51C can be well reproduced via p-p collisions of the accelerated protons in the part of shell interacting with the MC, and the radio emission from the SNR can be explained as the radiation of the electrons accelerated by the other part of the SNR shock freely expanding into the ambient interstellar space."
+ The pitch-anele averaged: steady-state. distribution of the protons accelerated at a shock in one dimension satisfies the dilfusive transport equation (e.g...Alalkov&Drury2001:Blasi2002:Amato.&Cabici2008) where the coordinate w is directed along the shock normal from downstream toupstream. 2 is the cillusion cocllicient and wis the [uid velocity in the shock frame. which equals 42. downstream. Gr.< 0). and. changes continuously upstream. from ij immediately upstream Gr=0 ) of the subshock to wy at upstream infinity Cr=| ox).," The pitch-angle averaged steady-state distribution of the protons accelerated at a shock in one dimension satisfies the diffusive transport equation \citep[e.g.,][]{MD01, B02, ABG08}, where the coordinate $x$ is directed along the shock normal from downstream toupstream, $D$ is the diffusion coefficient and $u$ is the fluid velocity in the shock frame, which equals $u_2$ downstream $x<0$ ) and changes continuously upstream, from $u_1$ immediately upstream $x=0^+$ ) of the subshock to $u_0$ at upstream infinity $x=+\infty$ )."
+ Lor the Bobm cdillusion. 2=pe(3cD). where D is the local magnetic field strength.," For the Bohm diffusion, $D=pc^2/(3eB)$, where $B$ is the local magnetic field strength."
+ With the assumption that the particles are injected at. immediate. upstream of the subshock. the source function can be written as Qu(p)o Gr).," With the assumption that the particles are injected at immediate upstream of the subshock, the source function can be written as $Q(x, p)=Q_0(p)\delta(x)$ ."
+ For monoenergetic injection. Qu(p) is where pi is the injection momentum. tas. ls the gas density ator =0 ancl 5g is the fraction of particles injected in the acceleration process.," For monoenergetic injection, $Q_0(p)$ is where $p_{\rm{inj}}$ is the injection momentum, $n_{\rm{gas, }1}$ is the gas density at $x=0^+$ and $\eta$ is the fraction of particles injected in the acceleration process."
+ With the injection recipe known as thermal leakage. 7 can be described. as gddhasXezἈπτήτ (Blasi.Gabici&Vannoni2005:Amato.Blasi&Cahier2008).. where Aa=f|us ds the compression factor at the subshock and £ is a parameter of theorder of 24 describing the injection momentum of the thermal particles in the downstream region (ping=£pi).," With the injection recipe known as thermal leakage, $\eta$ can be described as $\eta =
+4(R_{\rm{sub}}-1)\xi^3e^{-\xi^2}/3\pi^{1/2}$ \citep{BGV05, ABG08}, where $R_{\rm{sub}}=u_1/u_2$ is the compression factor at the subshock and $\xi$ is a parameter of theorder of 2–4 describing the injection momentum of the thermal particles in the downstream region $p_{\rm{inj}}=\xi p_{\rm{th,}2}$ )."
+" We use €=3.5 as in Amato.Blasi&Cabici(2008).. pusμμ.- is the thermal peak momentum of the particles in the downstream: [uid with temperature 15. my, is the proton mass and Ay is the Boltzmann constant."," We use $\xi=3.5$ as in \citet{ABG08}, $p_{\rm{th,}2}=(2m_pk_{\rm{B}}T_2)^{1/2}$ is the thermal peak momentum of the particles in the downstream fluid with temperature $T_2$ , $m_p$ is the proton mass and $k_{\rm{B}}$ is the Boltzmann constant."
+" Assuming the heating of the gas upstream is adiabatic. with the conservation condition of monicntum luxes between the two sides of the subshock. we can derive he relation between the temperature ofthe eas far upstream di and Z5. ie. =(54MàIa)HfHoalf(1suMS)Haas) to. where Ady is the Uuiel Mach number far upstream. £=Πο is the total compression actor. 5, is the ratio of specific heats (5,=5/3 for an ideal gas) Particles with larger momenta move farther away fron he shock than those with lower momenta. only particles with momentum 2p can reach the point or, (seedetailsinBlasi 2002).. and thus the pressure of the accelerated xwticles at the point can be described as where e(p) is the velocity of particles with momentum p."," Assuming the heating of the gas upstream is adiabatic, with the conservation condition of momentum fluxes between the two sides of the subshock, we can derive the relation between the temperature of the gas far upstream $T_0$ and $T_2$, i.e., $T_2 = (\gamma_g M_0^2/R_{\rm tot})[R_{\rm sub}/R_{\rm tot} -
+1/R_{\rm tot} + (1/\gamma_g M_0^2)(R_{\rm tot}/R_{\rm
+sub})^{\gamma_g}]T_0$ , where $M_0$ is the fluid Mach number far upstream, $R_{\rm tot}=u_0/u_2$ is the total compression factor, $\gamma_g$ is the ratio of specific heats $\gamma_g=5/3$ for an ideal gas) Particles with larger momenta move farther away from the shock than those with lower momenta, only particles with momentum $\geq p$ can reach the point $x_p$ \citep[see details in][]{B02}, , and thus the pressure of the accelerated particles at the point can be described as where $v(p)$ is the velocity of particles with momentum $p$."
+" Furthermore. the particle distribution function fu(p) at the shock can be implicith written as (Blasi2002) where may is the gas density far upstream (r=| ox). Duas is tho maximum momentum of the accelerated: protons. and it can be estimated as (Lagage&Cesarsky1983:Aha-ronian&Atovan1999) for a ""Bohm limit” according to the standard dilfusive shock acceleration. theory, where 6 the magnetic field strength. fs is the age of the host SNR in units of 105 ve."," Furthermore, the particle distribution function $f_0(p)$ at the shock can be implicitly written as \citep{B02}
+ where $n_{\rm ism}$ is the gas density far upstream $x=+\infty$ ), $p_{\rm max}$ is the maximum momentum of the accelerated protons, and it can be estimated as \citep{LC83,AA99}
+ for a ""Bohm limit"" according to the standard diffusive shock acceleration theory, where $B$ the magnetic field strength, $t_3$ is the age of the host SNR in units of $10^3$ yr."
+" ((p) can be solved. using an equation deduced. fron the conservation of the mass ancl momentum Iluxes with the boundary. condition (ping)=Bah and (puns)= 1. and then a value of 2,4, can be achieved. by an iterative procedure to satisfy the boundary conditions (Blasi 2002).."," $U(p)$ can be solved using an equation deduced from the conservation of the mass and momentum fluxes with the boundary condition $U(p_{\rm inj}) = R_{\rm
+sub}/R_{\rm tot}$ and $U(p_{\rm max})=1$ , and then a value of $R_{\rm sub}$ can be achieved by an iterative procedure to satisfy the boundary conditions \citep{B02}. ."
+ Theelectrons have the same spectrum of the protons up to a maximum energy determined by svnchrotron losses. and the spectrum can be described as (e.g..Fangetal.2009) where £(p) is the kinetic energy. of the electrons. and the electron/proton ratio Ao is treated.as a parameter.," Theelectrons have the same spectrum of the protons up to a maximum energy determined by synchrotron losses, and the spectrum can be described as \citep[e.g.,][]{Fea09}
+ where $E(p)$ is the kinetic energy of the electrons, and the electron/proton ratio $K_{ep}$ is treatedas a parameter."
+sample has a median value of aUV=—0.67 and the radio sample has a median value of aUV=—0.75.,sample has a median value of $\alpha^\mathrm{UV}_\nu=-0.67$ and the radio sample has a median value of $\alpha^\mathrm{UV}_\nu=-0.75$.
+" The lower panel of Figure 3 shows cumulative distributions of ffor the individual radio morphology classes, including the double-lobed quasars."," The lower panel of Figure \ref{fig:alphahists} shows cumulative distributions of for the individual radio morphology classes, including the double-lobed quasars."
+" For sources that have core radio emission, we apply a limit in radio flux density of 9.1 mJy (14.0mag,usingtheABmagnitudesystemof this is the limit above which the resolved/unresolved separation7); for the core sources is most effective (KI08)."," For sources that have core radio emission, we apply a limit in radio flux density of 9.1 mJy \citep[14.0 mag, using the AB magnitude system of][]{okeGunn}; this is the limit above which the resolved/unresolved separation for the core sources is most effective (KI08)."
+ 'The main result of this analysis is that unresolved core quasars have significantly higher fraction of extremely red sources thana all other quasar classes., The main result of this analysis is that unresolved core quasars have a significantly higher fraction of extremely red sources than all other quasar classes.
+" For example, unresolved cores make up of the radio sample with aUV<—1.5, but only of the radio sample with aLV>—0.5."," For example, unresolved cores make up of the radio sample with $\alpha^\mathrm{UV}_\nu<-1.5$, but only of the radio sample with $\alpha^\mathrm{UV}_\nu>-0.5$."
+" To quantify the statistical significance of differences in ddistributionsaV for radio quasars and radio-quiet quasars, we employ the Kolmogorov-Smirnov (KS) test."," To quantify the statistical significance of differences in distributions for radio quasars and radio-quiet quasars, we employ the Kolmogorov-Smirnov (KS) test."
+" The P-value resulting from such a test represents the probability of obtaining the observed distributions by chance, if the two samples share the same parent distribution."," The $P$ -value resulting from such a test represents the probability of obtaining the observed distributions by chance, if the two samples share the same parent distribution."
+ Table 4 shows the results of KS tests comparing the radio-quiet quasar distribution to each radio class., Table \ref{table:ks} shows the results of KS tests comparing the radio-quiet quasar distribution to each radio class.
+ The KS tests quantify what is shown qualitatively in Figure 3: the only quasar class with an extremely different (P<0.01) color distribution from the radio-quiet quasars is the unresolved core class., The KS tests quantify what is shown qualitatively in Figure \ref{fig:alphahists}: the only quasar class with an extremely different $P \ll 0.01$ ) color distribution from the radio-quiet quasars is the unresolved core class.
+" This difference must be a physical effect because unresolved core sources were targeted with the same algorithm as resolved cores, triples, lobes, and jets."," This difference must be a physical effect because unresolved core sources were targeted with the same algorithm as resolved cores, triples, lobes, and jets."
+" Because unresolved core quasars dominate the radio sample, a comparison of color distributions for the radio quasars and the radio-quiet quasars shows the difference reported by ?.."," Because unresolved core quasars dominate the radio sample, a comparison of color distributions for the radio quasars and the radio-quiet quasars shows the difference reported by \citet{i02}."
+ Double-lobed quasars were targeted using the same criteria as radio-quiet quasars., Double-lobed quasars were targeted using the same criteria as radio-quiet quasars.
+" Because those two classes have not dissimilar observed color distributions, it is possible that they were drawn from similar parent color distributions."," Because those two classes have not dissimilar observed color distributions, it is possible that they were drawn from similar parent color distributions."
+" The simplest interpretation of our results is: (1) unresolved core quasars are the only subsample of radio quasars that has an intrinsically redder color distribution than other quasars, and (2) optical colors of SDSS quasars are not strongly («0.1 mag) biased blue."," The simplest interpretation of our results is: (1) unresolved core quasars are the only subsample of radio quasars that has an intrinsically redder color distribution than other quasars, and (2) optical colors of SDSS quasars are not strongly $<0.1$ mag) biased blue."
+ We note that these results are specific to quasars in the redshift range zS2.7; quasars in Figure 3 are roughly limited to that range because of the redshift distribution of the S07 quasar catalog., We note that these results are specific to quasars in the redshift range $z\lesssim2.7$; quasars in Figure \ref{fig:alphahists} are roughly limited to that range because of the redshift distribution of the S07 quasar catalog.
+ It is not obvious why the unresolved cores have a much higher fraction of extremely red sources than the other samples., It is not obvious why the unresolved cores have a much higher fraction of extremely red sources than the other samples.
+" Figure 2 suggests that this result is not a redshift effect, because lobe and jet sources have similar redshift distributions to the unresolved core sample, while triples and resolved cores have similar redshift distributions to each other."," Figure \ref{fig:redshifts} suggests that this result is not a redshift effect, because lobe and jet sources have similar redshift distributions to the unresolved core sample, while triples and resolved cores have similar redshift distributions to each other."
+" As discussed in ?,, dust-reddened spectra are not easily characterized by a single power law."," As discussed in \citet{richards03}, dust-reddened spectra are not easily characterized by a single power law."
+ Our current analysis does not distinguish between spectra that are intrinsically red and those that are dust-reddened., Our current analysis does not distinguish between spectra that are intrinsically red and those that are dust-reddened.
+" We do not expect that the highly reddened sources in the unresolved core sample are caused by dust reddening, because these are sources that should have the smallest viewing angles and thus the least obscuration by the AGN torus."," We do not expect that the highly reddened sources in the unresolved core sample are caused by dust reddening, because these are sources that should have the smallest viewing angles and thus the least obscuration by the AGN torus."
+" Recent arguments suggest that iis a stronger statistical measure of core-boosting, and therefore of orientation, than R (see discussion in refsec:intro))."," Recent arguments suggest that is a stronger statistical measure of core-boosting, and therefore of orientation, than $R$ (see discussion in \\ref{sec:intro}) )."
+" R normalizes the core boosting using the extended radio emission, which depends on age and environment in addition to orientation."," $R$ normalizes the core boosting using the extended radio emission, which depends on age and environment in addition to orientation."
+" These factors should have less influence onRy,, which uses the core’s optical luminosity to normalize the core boosting."," These factors should have less influence on, which uses the core's optical luminosity to normalize the core boosting."
+" While 8 beamed synchrotron component (such as in blazars and BL Lac objects) would contaminate the eestimates in face-on sources, the presence of broad lines in the spectra for inclusion in the SDSS quasar sample) suggests that (requiredsuch a component is not seen in our sources."," While a beamed synchrotron component (such as in blazars and BL Lac objects) would contaminate the estimates in face-on sources, the presence of broad lines in the spectra (required for inclusion in the SDSS quasar sample) suggests that such a component is not seen in our sources."
+" In other words, it is unlikely that we have any directly face-on sources, as such orientations should result in an absence of any emission lines."," In other words, it is unlikely that we have any directly face-on sources, as such orientations should result in an absence of any emission lines."
+ We now compare values of R and, We now compare values of $R$ and
+VLBI observations only depend on the observing frequency.,VLBI observations only depend on the observing frequency.
+ The dynamie range ts thus The dynamic range of a phase-referenced image ts limited to a given value. Dj. This limit ts independent of both S and the sensitivity of the array.," The dynamic range is thus The dynamic range of a phase-referenced image is limited to a given value, $D_l$, This limit is independent of both $S$ and the sensitivity of the array."
+ It is achieved when the flux density of the target source is much higher than the thermal noise of the interferometer., It is achieved when the flux density of the target source is much higher than the thermal noise of the interferometer.
+ The limiting dynamic range only depends on @ v. and Ar. and it can be several orders of magnitude smaller than the corresponding dynamic range due to the thermal noise of the receiving system. indicating that the atmosphere strongly limits the sensitivity of the observations.," The limiting dynamic range only depends on $\theta$ , $\nu$, and $\Delta t$, and it can be several orders of magnitude smaller than the corresponding dynamic range due to the thermal noise of the receiving system, indicating that the atmosphere strongly limits the sensitivity of the observations."
+ Martí-Vidal et al. (2010)), Martí–Vidal et al. \cite{mar10}) )
+" propose a phenomenological model for fj, based on their Monte Carlo simulations (see their Sect.", propose a phenomenological model for $f_{at}$ based on their Monte Carlo simulations (see their Sect.
+ 5.3)., 5.3).
+ This model takes the form where KVAr is a constant to be determined., This model takes the form where $K$ is a constant to be determined.
+ We can compare this phenomenological model to our observations by plotting the dynamic range of all our phase-referenced images as a function of calibrator-to-target separation., We can compare this phenomenological model to our observations by plotting the dynamic range of all our phase-referenced images as a function of calibrator-to-target separation.
+ The images obtained for each source pair correspond to data with slightly different on-source observing times., The images obtained for each source pair correspond to data with slightly different on-source observing times.
+" Therefore. we corrected (1e.. we normalized) the dynamic ranges of all the images by applying the factor VNio/N,. where N, is the number of visibilities for each image and Nyy is the number of visibilities corresponding to a on-source observing time of 10 hours with the VLBA."," Therefore, we corrected (i.e., we normalized) the dynamic ranges of all the images by applying the factor $\sqrt{N_{10}/N_{v}}$ , where $N_{v}$ is the number of visibilities for each image and $N_{10}$ is the number of visibilities corresponding to a on-source observing time of 10 hours with the VLBA."
+ We show in Fig., We show in Fig.
+ | the normalized dynamic ranges and the model resulting from Eq. 3.., \ref{DynRPh} the normalized dynamic ranges and the model resulting from Eq. \ref{FAT}.
+ We find that for K~12.4hhr°> GGHz! the model predicts the limiting dynamic ranges obtained with the VLBA at both frequencies. although the results at GGHz are of higher quality: the results at GGHz are noisier.," We find that for $K \sim 12.4$ $^{0.5}$ $^{-1}$ the model predicts the limiting dynamic ranges obtained with the VLBA at both frequencies, although the results at GHz are of higher quality; the results at GHz are noisier."
+ The dynamic ranges (also normalized to an observing time of 10 hours) obtained from the self-calibrated images are shown in Fig., The dynamic ranges (also normalized to an observing time of 10 hours) obtained from the self-calibrated images are shown in Fig.
+ 2 for a comparison with those obtained from phase-referencing., \ref{DynRSf} for a comparison with those obtained from phase-referencing.
+ The error bars in Fig., The error bars in Fig.
+ | are set to be proportional to the flux density of the calibrator., \ref{DynRPh} are set to be proportional to the flux density of the calibrator.
+ This way the reader has information in asingle figure on the quality of the phase-referenced images and the quality of the calibrator visibilities., This way the reader has information in a single figure on the quality of the phase-referenced images and the quality of the calibrator visibilities.
+ We show in Fig., We show in Fig.
+ 3 the peak flux densities of the phase-referenced images relative to those obtained from the images corresponding to the self-calibrated visibilities as a function of distance to the calibrator., \ref{PEAKRATIOFIG} the peak flux densities of the phase-referenced images relative to those obtained from the images corresponding to the self-calibrated visibilities as a function of distance to the calibrator.
+ The systematics m the loss of flux density is clear for the GGHz data., The systematics in the loss of flux density is clear for the GHz data.
+ The flux density recovered is about of the real flux density for separations of ~5 degrees. and slowly decreases to for the largest separations.," The flux density recovered is about of the real flux density for separations of $\sim5$ degrees, and slowly decreases to for the largest separations."
+ For the GGHz data. the scatter is larger and no robust conclusion can be obtained.," For the GHz data, the scatter is larger and no robust conclusion can be obtained."
+ However. a hint of saturation of the peak ratios around can be appreciated at large source separations if the pointscorresponding to a separation of14.84 degrees (for which peak ratios as high as are obtained) are not considered (these points correspond," However, a hint of saturation of the peak ratios around can be appreciated at large source separations if the pointscorresponding to a separation of14.84 degrees (for which peak ratios as high as are obtained) are not considered (these points correspond"
+mean molecular weight and opacity in the lavers above the driving laver.,mean molecular weight and opacity in the layers above the driving layer.
+ The logical alternative is therefore to increase the iron abundance increasing the abundances ofother heavy elements., The logical alternative is therefore to increase the iron abundance increasing the abundances of other heavy elements.
+ Using the OPAL opacity table generator (Iglesias&Rogers 1996)... we constructed new sets of opacity tables in which the iron abundance alone was increased by [actors of f=5.10 and 20. whilst the remaining elements had abundances corresponding toa Z=0.02 mixture.," Using the OPAL opacity table generator \citep{Igl96}, we constructed new sets of opacity tables in which the iron abundance alone was increased by factors of $f = 5, 10$ and 20, whilst the remaining elements had abundances corresponding to a $Z=0.02$ mixture."
+ Phere is a modest effect on the overall Z such that Z—0.021(f.Zi. where Ze.=0.02188 represents the mass fraction of iron within the standard heavy. element mixture (Cirevesse&Noels 1993)., There is a modest effect on the overall $Z$ such that $Z=0.02+(f-1)Z_{\rm Fe}$ where $Z_{\rm Fe}=0.02188$ represents the mass fraction of iron within the standard heavy element mixture \citep{Gre93}.
+. A similar set of models. was run for these mixtures., A similar set of models was run for these mixtures.
+ These are compared. with the reference set 0.02 in Figs., These are compared with the reference set $Z=0.02$ in Figs.
+ 4 and 5.., \ref{Fe-stable} and \ref{Fe-modes}.
+ ]t is immediately. apparent that increasing the iron abundance alone by a factor five produces a predictable ellect similar to increasing Z from 0.02 to 0.10. but with the major difference that the blue edge has become bluer," It is immediately apparent that increasing the iron abundance alone by a factor five produces a predictable effect similar to increasing $Z$ from 0.02 to 0.10, but with the major difference that the blue edge has become bluer"
+size.,size.
+ Finally. Getman conclude that disk presence has no effect on peak flare luminosity and (otal flare energy (2008b). a result which we confirm in 844.," Finally, Getman conclude that disk presence has no effect on peak flare luminosity and total flare energy (2008b), a result which we confirm in 4."
+ The principle difference between our study and Gelman regards the role of disks in flaring behavior., The principle difference between our study and Getman regards the role of disks in flaring behavior.
+ As discussed in 855.1. we find that long flares of order Re occur only on disked PAIS stars anc conclude (hat in the case of these very long flares. one footpoint. of the magnetic loop must be anchored on the accretion disk.," As discussed in 5.1, we find that long flares of order $R_{C}$ occur only on disked PMS stars and conclude that in the case of these very long flares, one footpoint of the magnetic loop must be anchored on the accretion disk."
+ Gelman come to the opposite conclusion. observing Lares of order Hc: or longer almost exclusively on YSOs categorized as Class [I] in their study (2008b).," Getman come to the opposite conclusion, observing flares of order $R_{C}$ or longer almost exclusively on YSOs categorized as Class III in their study (2008b)."
+ While Getman find little actual variation in flare properties between clisked and diskless stars. they argue Chat due to (heir faster rotation periods. Class HI YSOs actually have significantly smaller corotation radii than Class II YSOs and so regardless of length in em. flares on Class I] YSOs rarely if ever reach the disk and are anchored on the stellar photosphere alone.," While Getman find little actual variation in flare properties between disked and diskless stars, they argue that due to their faster rotation periods, Class III YSOs actually have significantly smaller co–rotation radii than Class II YSOs and so regardless of length in cm, flares on Class II YSOs rarely if ever reach the disk and are anchored on the stellar photosphere alone."
+ Finally. we can compare our sample to the study of COUP Lares recently published by Aarnio (2010).," Finally, we can compare our sample to the study of COUP flares recently published by Aarnio (2010)."
+ Their goal was to understand the role that cireumstellar disks play in the enereelics of high-contrast. X-ray flares., Their goal was to understand the role that circumstellar disks play in the energetics of high-contrast X-ray flares.
+ The authors construct spectral energy distributions in the wavelength range 0.3 — 8 jun (extending in some cases out to 24 jan) and model them to determine whether (here is circumstellar disk material in sufficient proximity to the flares’ confining magnetic loops (o allow star-disk interaction., The authors construct spectral energy distributions in the wavelength range 0.3 $-$ 8 $\micron$ (extending in some cases out to 24 $\micron$ ) and model them to determine whether there is circumstellar disk material in sufficient proximity to the flares' confining magnetic loops to allow star-disk interaction.
+" IF 725. M, and 7;jj ave supplied. SED modelling can vield disk parameters including mass and sublimation radius."," If $R_{\star}$, $M_{\star}$ and $T_{eff}$ are supplied, SED modelling can yield disk parameters including mass and sublimation radius."
+ This technique characterizes (he circumstellar disk more quantitatively than the near-infrared color excesses used in (his study and has the advantage of detecting even cool disks with large inner holes., This technique characterizes the circumstellar disk more quantitatively than the near-infrared color excesses used in this study and has the advantage of detecting even cool disks with large inner holes.
+ Aarnio conclude that of the stars in their COUP high-contrast [lave sample have no disk material within reach of the confining magnetic loops and so argue that high-contrast X-ray [lares in general are purely stellar in origin. neither triggered nor stabilized by star-disk interactions.," Aarnio conclude that of the stars in their COUP high-contrast flare sample have no disk material within reach of the confining magnetic loops and so argue that high-contrast X-ray flares in general are purely stellar in origin, neither triggered nor stabilized by star-disk interactions."
+ The authors further argue (that as long as the confining magnetic field at the end of the loop is strong enough. the loop is stable even without being anchored on (he disk.," The authors further argue that as long as the confining magnetic field at the end of the loop is strong enough, the loop is stable even without being anchored on the disk."
+ We agree that disk presence has no effect on flare energetics. however our Figure 4. shows that confining magnetic field falls off sharply with loop leneth irrespective of disk presence and does not support this scenario.," We agree that disk presence has no effect on flare energetics, however our Figure \ref{BasicResults} shows that confining magnetic field falls off sharply with loop length irrespective of disk presence and does not support this scenario."
+ Aarnio do agree that for 10 of objects studied. star-disk interaction is possible and that the IR photometry of a further 8 sources in the sample does not permit conclusive statements about disks.," Aarnio do agree that for 10 of objects studied, star-disk interaction is possible and that the IR photometry of a further 8 sources in the sample does not permit conclusive statements about disks."
+ Of the 15 flares in their sample with loops longer than about 10 Πεν. 5 (lares appear within reach of the inner edee of the disk. ancl another 5 fall into the unknown category and have flare loop lengths of order the dust destruction radius.," Of the 15 flares in their sample with loops longer than about 10 $R_{\star}$, 5 flares appear within reach of the inner edge of the disk, and another 5 fall into the unknown category and have flare loop lengths of order the dust destruction radius."
+ While, While
+The discovery of tens of extra-solar planets over recent vears (Wolszezan 1994: Mavor Queloz 1995: Marcy Butler 2000) has re-invigorated efforts to understand the processes » which planetary svstems form.,The discovery of tens of extra-solar planets over recent years (Wolszczan 1994; Mayor Queloz 1995; Marcy Butler 2000) has re-invigorated efforts to understand the processes by which planetary systems form.
+ The existence of more han one svsten the solar svsteni to study and the oospect of many more being discovered: has spurrecl this elfort., The existence of more than one system — the solar system — to study and the prospect of many more being discovered has spurred this effort.
+ One wav to help understand the formation of planetary svstems ds. (o discover characteristics which distinguish planet-harbouring stars from. lone stars., One way to help understand the formation of planetary systems is to discover characteristics which distinguish planet-harbouring stars from lone stars.
+ They are more metal rich than the general stellar. population (Fubrmann. Pfeiller Bernkopf. L997: Gonzalez 1997. 1900). and the difference between solar photospheric ancl meteoric abundances correlates with elemental condensation temperature. consistent with self-enrichment of the solar surface. (Gonzalez 1997).," They are more metal rich than the general stellar population (Fuhrmann, Pfeiffer Bernkopf 1997; Gonzalez 1997, 1999), and the difference between solar photospheric and meteoritic abundances correlates with elemental condensation temperature, consistent with self-enrichment of the solar surface (Gonzalez 1997)."
+ Gonzalez (1999) discusses the anomalously small velocity of the sun relative to the local standard: of rest (LSIU). but the explanations are basec on anthropic arguments which do not tell us about other planetary svstems.," Gonzalez (1999) discusses the anomalously small velocity of the sun relative to the local standard of rest (LSR), but the explanations are based on anthropic arguments which do not tell us about other planetary systems."
+ Genuine characteristics not only provide information to help understand the formation of these systems. but could. also help bias future searches towards planet-harbouring systems of that type.," Genuine characteristics not only provide information to help understand the formation of these systems, but could also help bias future searches towards planet-harbouring systems of that type."
+ Perusal of the characteristics of cxoplanct hosts can eive the impression that they are unusually Li deficient compared to lone stars., Perusal of the characteristics of exoplanet hosts can give the impression that they are unusually Li deficient compared to lone stars.
+ Several stars now known to have xdanetary systems were Πασσος as having low Li abundancesprior to the discovery of their companions., Several stars now known to have planetary systems were flagged as having low Li abundances to the discovery of their companions.
+ LR 5968 (p CrD) was singled out. by Lambert. Heath. Eclvardsson 991). and Priel et al. (," HR 5968 $\rho$ CrB) was singled out by Lambert, Heath Edvardsson (1991), and Friel et al. ("
+1993) commented on the large Li difference between 16 €ve A and. B despite their similar emperatures. though they did not suggest that processes (her than normal single-star evolution would. be needed o explain the lower abundance in 16 Cvg D. As a third example. the low Li abundance in the solar photosphere (.MLi) = L.10+0.10: CGrevesse Sauval 1998). compared with the pre-solar nebula (CA(Li) = 3.31—0.04 in meteorites) was long challenged standard stellar evolution moclels (c.g. Delivannis 1995).,"1993) commented on the large Li difference between 16 Cyg A and B despite their similar temperatures, though they did not suggest that processes other than normal single-star evolution would be needed to explain the lower abundance in 16 Cyg B. As a third example, the low Li abundance in the solar photosphere $A$ (Li) = $\pm$ 0.10; Grevesse Sauval 1998) compared with the pre-solar nebula $A$ (Li) = $\pm$ 0.04 in meteorites) has long challenged standard stellar evolution models (e.g. Deliyannis 1995)."
+ CALI) = logio (»(Li)/o(11)) |. 12.00.), $A$ (Li) $\equiv$ $_{10}$ $n$ $n$ (H)) + 12.00.)
+ Lithium is special because stars destroy it during pre-main sequence and. main-sequence evolution. depending on their mass ane metallicity.," Lithium is special because stars destroy it during pre-main sequence and main-sequence evolution, depending on their mass and metallicity."
+ When surface material is mixed down to depths where the temperature exceeds 10° ky. Li-purged material is returned to the surface.," When surface material is mixed down to depths where the temperature exceeds $\times$ $^6$ K, Li-purged material is returned to the surface."
+ Li survival therefore rellects the mixing history. and in the context of planet-harbouring stars could. provide information on the accretion of material and the angular-momoentum evolution of the system as a whole.," Li survival therefore reflects the mixing history, and in the context of planet-harbouring stars could provide information on the accretion of material and the angular-momentum evolution of the system as a whole."
+ Li deficiency in planet hosts was assessed by Ixing ct al. , Li deficiency in planet hosts was assessed by King et al. (
+997) and Gonzalez Laws (2000).,1997) and Gonzalez Laws (2000).
+ Ixing et al., King et al.
+ examined 16 Cye A and D. and commented. on six other systems.," examined 16 Cyg A and B, and commented on six other systems."
+ LID 114762. 70 Vir. and 7 Boo.," HD 114762, 70 Vir, and $\tau$ Boo."
+" They concluded that ""the data are too few at this point to establish a connection", They concluded that “the data are too few at this point to establish a connection
+other hand. alternate accretion outbursts that typically [ast [or weeks to months with vears to decades long episodes of cquiescence. during which the X-ray. luminosity is more than 2 orders of magnitude lower.,"other hand, alternate accretion outbursts that typically last for weeks to months with years to decades long episodes of quiescence, during which the X-ray luminosity is more than 2 orders of magnitude lower."
+ One of the phenomena that uniquely mark the compact primary as à neutron star are type-L N-ray bursts (or shortly “X-ray bursts’): bright Πάθος of X-rav emission that are caused by unstable nuclear burning on the surface of the neutron star., One of the phenomena that uniquely mark the compact primary as a neutron star are type-I X-ray bursts (or shortly `X-ray bursts'); bright flashes of X-ray emission that are caused by unstable nuclear burning on the surface of the neutron star.
+" They are characterized by blackbody emission with a peak temperature Aia,>2 keV and generally display a last rise time followed: by a slower decay phase.", They are characterized by blackbody emission with a peak temperature $kT_{\mathrm{bb}}>2$ keV and generally display a fast rise time followed by a slower decay phase.
+ The initial rise can be interpreted: as burning of the fuel [aver. while the subsequent decay. represents the cooling of the ashes.," The initial rise can be interpreted as burning of the fuel layer, while the subsequent decay represents the cooling of the ashes."
+ So far. X-ray bursts have only been detected. from low-mass X-ray binaries (ΤΕΛΙΝBs). in which the donor star has a mass Al1M.," So far, X-ray bursts have only been detected from low-mass X-ray binaries (LMXBs), in which the donor star has a mass $M \lesssim 1~\Msun$."
+ Lhe properties (e.g.. duration. radiated cnerey and recurrence-. time)) of type-b X-ray bursts depend on the conditions of the ignition laver. such as the temperature. thickness and hydrogen (11) abundance.," The properties (e.g., duration, radiated energy and recurrence time) of type-I X-ray bursts depend on the conditions of the ignition layer, such as the temperature, thickness and hydrogen (H) abundance."
+ These can drastically change as the mass-acerction rate onto the neutron star varies. which results in X-ray bursts with cillerent characteristies for dillerent accretion regimes (Lorreviews.seee.g...Lewin.vanParadijs&Taam1995:Strohmaver&Bildsten 2006).," These can drastically change as the mass-accretion rate onto the neutron star varies, which results in X-ray bursts with different characteristics for different accretion regimes \citep[for reviews, see e.g., ][]{lewin95,strohmayer06}."
+. X-rav bursts can be serendipitously. detected. by the Durst Alert Telescope (BAT:Barthelmyetal.2005). aboard he ssalellite: a multi-wavelength observatory that is dedicated o the study of gamma-ray bursts (CGl1tDs)., X-ray bursts can be serendipitously detected by the Burst Alert Telescope \citep[BAT;][]{barthelmy05} aboard the satellite; a multi-wavelength observatory that is dedicated to the study of gamma-ray bursts (GRBs).
+ Although events rom known X-rav burst sources are ignored. the BAT occasionally triggers on an X-ray burst. from a previously unknown burster (e.g.intZandetal.2008:Linaresοἱ2009:Wijnanelsctal. 2009).," Although events from known X-ray burst sources are ignored, the BAT occasionally triggers on an X-ray burst from a previously unknown burster \citep[e.g.,][]{zand08,linares09,wijnands09}."
+. On 2008 May. 14 at 10:32:37 WE. Swiff'ss BAT registered. an X-ray Uare (Ixrimm.et 2008).," On 2008 May 14 at 10:32:37 , s BAT registered an X-ray flare \citep[][]{krimm08}."
+. The BAT lDighteurve ancl soft. X-ray spectrum (no photons detected above ~35 keV) suggested that this event was not a GiB (Daumgartneoret.al.2008:Ixrimmetal. 2008).," The BAT lightcurve and soft X-ray spectrum (no photons detected above $\sim35$ keV) suggested that this event was not a GRB \citep{baumgartner08,krimm08}."
+. Rapid follow-up observations with the X-ray Telescope (XII:Burrowsetal.2005) detected a bright. rut quickly fading A-ray source within the 3 aremin BAP error circle (Ixrimmetal.200S:Daumgartner2008)," Rapid follow-up observations with the X-ray Telescope \citep[XRT;][]{burrows05} detected a bright, but quickly fading X-ray source within the 3 arcmin BAT error circle \citep{krimm08,baumgartner08}."
+ simultaneously. obtained UV/Optical Telescope (UVOT:toming2005) images revealec a [ading optical source within the XICE error circle (Israeletal.2008) , Simultaneously obtained UV/Optical Telescope \citep[UVOT;][]{roming05} images revealed a fading optical source within the XRT error circle \citep{israel08}. .
+"The UVOT detection allowed for an accurate ocalization of the source. of the MVE grigeecr: ao—Ίταπο της, ὃ=v4016.1"". (2000) wit a per cent confidence radius of{""OG arescc nmetal.withhen2008)."," The UVOT detection allowed for an accurate localization of the source of the BAT trigger: $\alpha= \mathrm{17^{h}35^{m}23.75^{s}}$ , $\delta=-35^{\circ} 40' 16.1\arcsec$ (J2000) with a 90 per cent confidence radius of 0.56 arcsec \citep{israel08}."
+. Both he ARP and the UV position coincide that of the unclassified X-ray sourcerxu hereafter). ich was discovered. with the ssatellite in. 1990," Both the XRT and the UVOT position coincide with that of the unclassified X-ray source; hereafter), which was discovered with the satellite in 1990."
+ The BAT trigger was Likely the result of an X-ray burst [rom this svstem (Israeletal.2008)... and would thereby identify aas a neutron star in. most likely. an LMXD.," The BAT trigger was likely the result of an X-ray burst from this system \citep{israel08}, and would thereby identify as a neutron star in, most likely, an LMXB."
+ We note that IHtodriguezetal.(2009). useπο XR observations discussed in this paper to obtain a 3.5 aresec position for the likely hard. X-ray counterpart ofJ1735.. ((sce Section 27)).," We note that \citet{rodriguez09} used /XRT observations discussed in this paper to obtain a 3.5 arcsec position for the likely hard X-ray counterpart of, (see Section \ref{subsec:integral}) )."
+ Dased on that position. the authors identify a brighto counterpart candidate in PALASS (4¥s==8.63+0.03) and USNO-B1L.O —2—11.9+0.3) catalogues. sugeesting a possible high-mass X-ray. binary nature.," Based on that position, the authors identify a bright counterpart candidate in 2MASS $8.63\pm0.03$ ) and USNO-B1.0 $11.9\pm0.3$ ) catalogues, suggesting a possible high-mass X-ray binary nature."
+ μις object is also visible in our optical anc near-infrared (near-LR) observations. but although it is very close to ((~+4 arcsec NW: sce Fig. 3)).," This object is also visible in our optical and near-infrared (near-IR) observations, but although it is very close to $\sim4$ arcsec NW; see Fig. \ref{fig:images}) ),"
+ it lies well outside the sub-aresecond UVOT. position and is therefore not. its counterpart., it lies well outside the sub-arcsecond UVOT position and is therefore not its counterpart.
+ In this paper we report on a multi-wavelength observing campaign of {following the BAT trigger of 2008 Alay 1H., In this paper we report on a multi-wavelength observing campaign of following the BAT trigger of 2008 May 14.
+ We discuss the properties of the X-ray burst ancl the characteristics of the persistent emission., We discuss the properties of the X-ray burst and the characteristics of the persistent emission.
+ Our study comprises delata obtained with the BAT. NIE. and UVOT. optical photometric observations carried. out with the (REAL)) and the (NTTL)). optical spectroscopy using the (VLT)). as well as near-LH observations performed with the telescope.," Our study comprises data obtained with the BAT, XRT and UVOT, optical photometric observations carried out with the ) and the ), optical spectroscopy using the ), as well as near-IR observations performed with the telescope."
+ In addition. we explore archivalROSAT.. aancl cedata to investigate the long-term. Εις and X-ray burst behavior of1735.," In addition, we explore archival, and data to investigate the long-term flux and X-ray burst behavior of."
+. The observations that we obtained of wwith cülferent facilities are. listed in ‘Table 1.., The observations that we obtained of with different facilities are listed in Table \ref{tab:obs}.
+ 1n the following sections these are discussed in more detail., In the following sections these are discussed in more detail.
+ We generated standard. BAT data products for the trigger observation using the tool., We generated standard BAT data products for the trigger observation using the tool.
+ Phe 1535 keV BAT lighteurve of the burst. shown in Fig. 1.," The 15–35 keV BAT lightcurve of the burst, shown in Fig. \ref{fig:bat},"
+ ls consistent with a single peak centred at /— Os and emerging [rom the [or 200 s. with a very slow rise time of 100 s --rd(Baumgartnerctal.2008:Israelet20," is consistent with a single peak centred at $t\sim0$ s and emerging from the background for $\sim 200$ s, with a very slow rise time of $\sim 100$ s \citep{baumgartner08,israel08}."
+08)DA... In Fig. 1.," In Fig. \ref{fig:bat},"
+ the apparent peak at /~90 s is thought to be an artefact related to the spacecraft. slewing. whereas the apparent rise in count rate after /~120 is likely caused by entering the South Atlantic Anomaly (SAA:Baumeartneretal. 2008).," the apparent peak at $t\sim90$ s is thought to be an artefact related to the spacecraft slewing, whereas the apparent rise in count rate after $t\sim120$ is likely caused by entering the South Atlantic Anomaly \citep[SAA;][]{baumgartner08}."
+. The spacecraft started slewing 75 s after the burst trigeer. by which time the BAT count rate had nearly dropped. to the background. level(see. Fig.," The spacecraft started slewing $\sim75$ s after the burst trigger, by which time the BAT count rate had nearly dropped to the background level (see Fig. \ref{fig:bat}) )."
+ Therefore. we used only pre-slew cata ancl extracted: a single BAT spectrum. of 140 s around the burst peak using the toolDBAFBINEVT.," Therefore, we used only pre-slew data and extracted a single BAT spectrum of 140 s around the burst peak using the tool."
+. Cüven the low count rate. it is notuseful to divide the BAT data in multiple bins with a higher time resolution.," Given the low count rate, it is notuseful to divide the BAT data in multiple bins with a higher time resolution."
+ Necessary geometrical corrections were applied with and the BDAT-recommenced svstematical error was administered using BATPUASYSERR., Necessary geometrical corrections were applied with and the BAT-recommended systematical error was administered using .
+. We generated asingle response matrix bv running the task, We generated asingle response matrix by running the task
+ou he cwnamiucal history of the solar svstei by requiring that Neptune not disrupt these objects as it migrates outward on au 1wclined and/or eccentric orbit.,on the dynamical history of the solar system by requiring that Neptune not disrupt these objects as it migrates outward on an inclined and/or eccentric orbit.
+" In this SCClon. we discuss our specific criteria for ""preserving the cold classicals. which we will use to place constraints O11 he set of paralcters (Section 2.1)) defining Neptune's orbital historv."," In this section, we discuss our specific criteria for “preserving"" the cold classicals, which we will use to place constraints on the set of parameters (Section \ref{subsec:param}) ) defining Neptune's orbital history."
+ There is evidence that the classical KBOs have a ποσα inclination distribution (???)..," There is evidence that the classical KBOs have a bimodal inclination distribution \citep{2001B,2010G,2011V}."
+ The cold Classicals are defined as the class of objects with a distribution of iuclinatious / ceitered on a low inclination with a small width iu iucliuaion., The cold classicals are defined as the class of objects with a distribution of inclinations $i$ centered on a low inclination with a small width in inclination.
+ The functional form O lis distribution is typically modeled as a Camssian nuutiplied bv siu/., The functional form of this distribution is typically modeled as a Gaussian multiplied by $\sin i$.
+ One (7?) of the three proposed 1Clels for the de-biased inclination distribution differs substantially from the other two (22). with respect to the relative populations of the cold aud hot classicals aud the width of the hot (high /) component (Fie. 1..," One \citep{2010G} of the three proposed models for the de-biased inclination distribution differs substantially from the other two \citep{2001B,2011V} with respect to the relative populations of the cold and hot classicals and the width of the hot (high $i$ ) component (Fig. \ref{fig:distcompare},"
+ top panel)., top panel).
+ Ilowever. all three distributions are similar for the cold classicas (Fig. l..," However, all three distributions are similar for the cold classicals (Fig. \ref{fig:distcompare},"
+ bottom panel)., bottom panel).
+ The number of cold classicas per inclination biu falls off alinost cutirely by /=G for the models of ? and ο and bv;=LE for the model of ?.., The number of cold classicals per inclination bin falls off almost entirely by $i = 6^\circ$ for the models of \citet{2001B} and \citet{2010G} and by $i = 4^\circ$ for the model of \citet{2011V}.
+ Therefore we require that Neptune's dynamical history should not excite the cold classicals above au inclination of 6 degrees., Therefore we require that Neptune's dynamical history should not excite the cold classicals above an inclination of 6 degrees.
+ The disruption criterion for the cold classical KBO eceelyvicitics is more subtle., The disruption criterion for the cold classical KBO eccentricities is more subtle.
+ The cold population is defined by its iclination distribution. uot its eccentricity distrinition.," The cold population is defined by its inclination distribution, not its eccentricity distribution."
+ We could imagine a cold population of objecs which have melinatious below six degrees but a uniform distribution of eccentricities., We could imagine a cold population of objects which have inclinations below six degrees but a uniform distribution of eccentricities.
+ If this were true (ale Lowe will demonstrate that it is not). the initial planeesinal disk could be excited to arbitrarily large ecceutricities.," If this were true (and we will demonstrate that it is not), the initial planetesimal disk could be excited to arbitrarily large eccentricities."
+ Moreover. the eccentricity distribution could be shaped entirely by the lone-term stability of the KDBO orbits.," Moreover, the eccentricity distribution could be shaped entirely by the long-term stability of the KBO orbits."
+ In this case. objects excited in eccentrieitv durus Neptune's liel-eccentricity period would be ejected from the system over billions of vears.," In this case, objects excited in eccentricity during Neptune's high-eccentricity period would be ejected from the system over billions of years."
+ To test whether the eccentricities of the cold classicals are sculpted solely by stabiitv. we compared observed cold classical objects to as ability map created by ? (Fig. 2)).," To test whether the eccentricities of the cold classicals are sculpted solely by stability, we compared observed cold classical objects to a stability map created by \citet{2005L} (Fig. \ref{fig:diagnostic}) )."
+ 2? do not use proovr Clements. so we use the instantaneous orbital elements of the observed objects.," \citet{2005L} do not use proper elements, so we use the instantaneous orbital elements of the observed objects."
+ We have confirmed that the eatures of the distributious we ideutifv below are qualitaively the sane using proper elements (2.AppendixA.1»., We have confirmed that the features of the distributions we identify below are qualitatively the same using proper elements \citep[][Appendix A.1]{2012D}.
+ Because the hot aud cold population overlap (Fig. 1))," Because the hot and cold population overlap (Fig. \ref{fig:distcompare}) ),"
+ we cannot definitively determine to which distribution any particular object belongs., we cannot definitively determine to which distribution any particular object belongs.
+" However. for all three model distributions. less than 10 of objects with Πιο]nations /«2"" are hot."," However, for all three model distributions, less than $\%$ of objects with inclinations $i<2^\circ$ are hot."
+ We find that the ecceutricitics of /«2? objects in the region from 12.5-15 AU are confined wel below the survival limit., We find that the eccentricities of $i<2^\circ$ objects in the region from 42.5-45 AU are confined well below the survival limit.
+ Consequenutlv. we can οςunservativelv coustraiu the dynamical history of Neptune: Neptune cannot excite the cold classical objects iu this τε'elon above e 0.1.," Consequently, we can conservatively constrain the dynamical history of Neptune: Neptune cannot excite the cold classical objects in this region above e = 0.1."
+ Thus we can impose two conservative criteria for preserving the cold classicals: Excitation of the planetesinal population by a perturbing planet on an inclined and/or eccentric orbit occurs through secular evolution., Thus we can impose two conservative criteria for preserving the cold classicals: Excitation of the planetesimal population by a perturbing planet on an inclined and/or eccentric orbit occurs through secular evolution.
+ Here we preseut simple analytical expressions that we will use to predict ancl interpret the results of our iutcerations im Section L.., Here we present simple analytical expressions that we will use to predict and interpret the results of our integrations in Section \ref{sec:results}.
+ Due to forcing from. both the ecceutricity and inclination of a planctesimal nudereo secular oscillations on timescales of οτίer a nüllion vears.," Due to forcing from, both the eccentricity and inclination of a planetesimal undergo secular oscillations on timescales of order a million years."
+ The auetesiaals total eccentricity is the vector stun ofits Orced eccenVICHY 6€gaesgs Bmuaparted by Neptine. and its LEE CCCOLTLICITY Cy. set bv nlial conditions (Fig. 3)).," The planetesimal's total eccentricity is the vector sum ofits forced eccentricity $\eforced$, imparted by Neptune, and its free eccentricity $\efree$, set by initial conditions (Fig. \ref{fig:forcefree}) )."
+ The ομως vector precesses about the ege vector at the aneular frecποιος gkpno., The $\efree$ vector precesses about the $\eforced$ vector at the angular frequency $\gkbo$.
+ The secular evoluion of the Ροesta‘s duclination is aua ogousfo vet. to lowest order. separable frou the eccentricity evoution.," The secular evolution of the planetesimal's inclination is analogous to – yet, to lowest order, separable from – the eccentricity evolution."
+ The ylanctesimal’s total inclin:dion is the sun of its forced and free inclination. aud he free inclination precesses about the forced inclinatiπα.," The planetesimal's total inclination is the sum of its forced and free inclination, and the free inclination precesses about the forced inclination."
+" Iu the case of iuclination. we can think of the particle""s orbit being inclined by po. with respect to the ""force7 plane and precessing about the forced plane at the rate! gkpo."," In the case of inclination, we can think of the particle's orbit being inclined by $\ifree$ with respect to the “forced"" plane and precessing about the forced plane at the rate $\gkbo$."
+ sce ?.. Chapter 7. for a pedagogical prescutation of secular evolution.," See \citet{2000M}, , Chapter 7, for a pedagogical presentation of secular evolution."
+ The vector coniponeis οἳ the. plauctesimal’s eccentricity are |7οιcx aud fk=e cosa. where = is the planetesimal's longitue ofperiapse.," The vector components of the planetesimal's eccentricity are $h = e \sin \varpi$ and $k = e \cos \varpi$ , where $\varpi$ is the planetesimal's longitude ofperiapse."
+ Secular forcing by Neptune causes / aud & to evolve as (to first order iu e aud ον):, Secular forcing by Neptune causes $h$ and $k$ to evolve as (to first order in $e$ and $e_N$ ):
+One well-known siortcoming: of tle APLC is its chromaicity.,One well-known shortcoming of the APLC is its chromaticity.
+ This could be compensated using standard linear ort—-nization techüLiques {ο Create an apocizatiou that concentrates light and uiininizes sidelobes ac‘oss a desiguatec bandpass., This could be compensated using standard linear optimization techniques to create an apodization that concentrates light and minimizes sidelobes across a designated bandpass.
+ Alternativev. the mask could simply be sized for the longest. waveleiLet1 iun the baid. while shorter waveleiigths create narrower PSEs whose cores are still blocked bN7 the mask.," Alternatively, the mask could simply be sized for the longest wavelength in the band, while shorter wavelengths create narrower PSFs whose cores are still blocked by the mask."
+ Tus may still scatter Leh in uudesirable ways in the image plane., This may still scatter light in undesirable ways in the image plane.
+ A third option is tte use of eligelivalues less than Ag. wlucho may suppress more light off the central waveleneth while still inaintainiue high suppression.," A third option is the use of eigenvalues less than $\Lambda_0$, which may suppress more light off the central wavelength while still maintaining high suppression."
+ This metiod is employed in the APLC desigu for CPI (?).., This method is employed in the APLC design for GPI \mycitep{Mac08}.
+ Some investigation would be required to cletermine which method is preferable. depeudiug on the telescope arraugeime," Some investigation would be required to determine which method is preferable, depending on the telescope arrangement."
+ Te majority of large telescopes today have obstructed apertures auc alt-az mounts., The majority of large telescopes today have obstructed apertures and alt-az mounts.
+ Field ο will cause the binary pair to rotate i1 the image plane. with each star maintaining a copy of the teles‘ope PSF.," Field rotation will cause the binary pair to rotate in the image plane, with each star maintaining a copy of the telescope PSF."
+ The'é are tvνο appro:iclies that can be take to compensate for the field ‘olatlou: siujiulta160115 ««)unter-rolon ol the pupil and Lyot st«yp with the image rotator activated. aud rotatiol of tle inuage-plaue 1lasks with the image rotator deactivaed.," There are two approaches that can be taken to compensate for the field rotation: simultaneous counter-rotation of the pupil and Lyot stop with the image rotator activated, and rotation of the image-plane masks with the image rotator deactivated."
+" Ppil rotation has been ested prevνε, ou sky: the ptpil-trackine moe of the NACO instrujent o the VLT rotates the entire instrumen to keep the pup fixed (??).."," Pupil rotation has been tested previously on sky: the pupil-tracking mode of the NACO instrument on the VLT rotates the entire instrument to keep the pupil fixed \mycitep{Kas09a, Tut10}."
+ While eitley is possible. rotatine the image-plane uasks may be preferred. as it nmilunizes 11e nuuber of actUuators reetired: the image-plane masks uust be actuated regatless. to match the separation of tle target binary.," While either is possible, rotating the image-plane masks may be preferred, as it minimizes the number of actuators required; the image-plane masks must be actuated regardless, to match the separation of the target binary."
+ ΤΙe simplest way to place tle Wes ksin the imagee plate is with open-loop coutrol-kleutify the image plane locations fo ‘the masks wihout the masks in. then move them iuto position aid make science observations.," The simplest way to place the masks in the image plane is with open-loop control–identify the image plane locations for the masks without the masks in, then move them into position and make science observations."
+ A fast steejug nΗΤΟΙ may be used o provide coarse correction. though as here are two sources beiug blocked iuxepedently. at least one mask will still have to be adjusted.," A fast steering mirror may be used to provide coarse correction, though as there are two sources being blocked independently, at least one mask will still have to be adjusted."
+ If this cau be aligued wihi sullicient precision. the alieuimeit procedure ends here.," If this can be aligned with sufficient precision, the alignment procedure ends here."
+ If necessary. we sug[n]€rest an ackditioial. closed-loop control t«) 1nalutaln mask alignment.," If necessary, we suggest an additional closed-loop control to maintain mask alignment."
+ control will also be jecessary iasvstem is run with the iuiage rotator oll. to eusure the masks rotate with the ilage plane.," Closed-loop control will also be necessary if a system is run with the image rotator off, to ensure the masks rotate with the image plane."
+ One possible method of performiig closed-loop control is to use flux in the imeige plane as a feeback signal., One possible method of performing closed-loop control is to use flux in the image plane as a feedback signal.
+ Decenteriug the masks wilI tend to icrease [Iux sharply. as this is equivalent ο Inticxluciug ilt errors. which increases the fhx leaking twough an APLC (?)..," Decentering the masks will tend to increase flux sharply, as this is equivalent to introducing tilt errors, which increases the flux leaking through an APLC \mycitep{Siv08}."
+ We note that poiling cont‘ol using scieice Camera imagery lias alreacly beet demonstrated ou-scy with a vector vorex Inask. in the imaging of HR&799 from Palomar (? 26though the tecliniqe used to close the OOD Was 100 the same.," We note that pointing control using science camera imagery has already been demonstrated on-sky with a vector vortex mask, in the imaging of HR8799 from Palomar \mycitep{Ser10}, although the technique used to close the loop was not the same."
+ The total energy in tle pudil plaie ls shown in for a number of sky angles iu t vicinity of a mask., The total energy in the pupil plane is shown in for a number of sky angles in the vicinity of a mask.
+ As sheyw. the regiou ¢X suppression surrounding the co‘rect alignimeut is clistine though smearecl by atiospheric errors.," As shown, the region of suppression surrounding the correct alignment is distinct, though smeared by atmospheric errors."
+ Iuproved AO correction will improve the sensitivity., Improved AO correction will improve the sensitivity.
+ Movi180 downward into this region would be relatively straightforward. if oue mask is hekd fixed.," Moving downward into this region would be relatively straightforward, if one mask is held fixed."
+ This has, This has
+complication arises from the presence of the large-scale low entropy gas associated with the infalling group associated with the NGC 4911 galaxy. first detected by White. Briel and Henry (1993) and further studied by Vikhlinin et al. (,"complication arises from the presence of the large-scale low entropy gas associated with the infalling group associated with the NGC 4911 galaxy, first detected by White, Briel and Henry (1993) and further studied by Vikhlinin et al. ("
+1997) and Neumann et al. (,1997) and Neumann et al. (
+2003).,2003).
+ The group occupies the south-eastern part of Fig.6.., The group occupies the south-eastern part of \ref{f:imhent}.
+ A number of weaker large-scale features are also evident in Fig.6.., A number of weaker large-scale features are also evident in \ref{f:imhent}.
+ To summarize. we decided to use the following identification criterion: if a negative valley in the entropy map or a positive peak in the emission map lies within the size of the galaxy. it is considered to be associated with the galaxy.," To summarize, we decided to use the following identification criterion: if a negative valley in the entropy map or a positive peak in the emission map lies within the size of the galaxy, it is considered to be associated with the galaxy."
+ The resulting detections of galaxies are presented in Tab.1 and Tab.2 for the emission and entropy methods. respectively.," The resulting detections of galaxies are presented in \ref{t:x} and \ref{t:s} for the emission and entropy methods, respectively."
+ The entropy method is only reported for the center of the Coma cluster where the level of the overall emission is high., The entropy method is only reported for the center of the Coma cluster where the level of the overall emission is high.
+ As discussed below. most of the point sources detected in the present survey are identified with backgrourd AGNs.," As discussed below, most of the point sources detected in the present survey are identified with background AGNs."
+ In order to compare with other AGN studies. we selected the 2.0 and 2.0—4.5 keV bands.," In order to compare with other AGN studies, we selected the 0.5--2.0 and 2.0–4.5 keV bands."
+ The flux is extracted tsing a circle of 20” radius. centered on the detected peak.," The flux is extracted using a circle of $20^{\prime\prime}$ radius, centered on the detected peak."
+ The correction for flux outside these apertures is 1/0.74 in the 0.5—2.0 keV band and 1/0.73 in the 2.0—4.5 keV band. using the results of the in- calibrations of Ghizzardi (2001).," The correction for flux outside these apertures is 1/0.74 in the $0.5-2.0$ keV band and 1/0.73 in the $2.0-4.5$ keV band, using the results of the in-flight calibrations of Ghizzardi (2001)."
+ Changes in the PSF at large off-axis angles lead to variation of the aperture correction within for both the soft and the hard band., Changes in the PSF at large off-axis angles lead to variation of the aperture correction within for both the soft and the hard band.
+ The countrate-to-flux conversion for 0.5—2.0 keV band was done assuming a power law spectrum with a photon index of 1.4. which for the medium filter used in this survey. gives an effective area of the telescope of 1057 em? and a mean energy of 1.03 keV.," The countrate-to-flux conversion for $0.5-2.0$ keV band was done assuming a power law spectrum with a photon index of 1.4, which for the medium filter used in this survey, gives an effective area of the telescope of 1057 $^2$ and a mean energy of 1.03 keV."
+are enhanced ta starburst reeious. ds also and would explain the ITI] A5007/1II3 ratio of~ —3 pose(Storchi-Beremann et al. 1995)).,"are enhanced in starburst regions, is also possible and would explain the ] $\lambda$ $\beta$ ratio of $\sim$ 3 (Storchi-Bergmann et al. \cite{storchi-kinney}) ),"
+ lower than observed in ACNs., lower than usually observed in AGNs.
+ According to the inified: model CAutonucci 1993)). broad cuussion lines are extinguishedc» im Sevfor 2 ealaxics bv a dusty torus located at 100 pe from the uucleus.but they should be scattered iuto our line of sigh by free electrous ivouud he torus aud should be observed in poluized liebt.," According to the `unified model' (Antonucci \cite{antonucci}) ), broad emission lines are extinguished in Seyfert 2 galaxies by a dusty torus located at $\le$ 100 pc from the nucleus, but they should be scattered into our line of sight by free electrons around the torus and should be observed in polarized light."
+ Towever. Ueisler et al. (19973) ," However, Heisler et al. \cite{heisler}) )"
+did no detect broad cussion lines in the polarized specuiu of NGC 7582: they attributed the non-detection of polarize broad lines to the fact that our line of sight passes through an amount of dust in the torus which is lavecr than iu Sevtert 28 with polarized broad lines., did not detect broad emission lines in the polarized spectrum of NGC 7582: they attributed the non-detection of polarized broad lines to the fact that our line of sight passes through an amount of dust in the torus which is larger than in Seyfert 2s with polarized broad lines.
+ The presence of strong star formation activity in the ceutral kpe was concluded from other indicators., The presence of strong star formation activity in the central kpc was concluded from other indicators.
+ Morris et al. (1985)), Morris et al. \cite{morris}) )
+ showed that iu the inner Isx: the Πα cussion suggests a rotating disk of HII regions in the plane of he galaxy., showed that in the inner kpc the $\alpha$ emission suggests a rotating disk of HII regions in the plane of the galaxy.
+ Recently. Arctxaga et al. (1999)), Recently Aretxaga et al. \cite{aretxaga}) )
+ reported the appearauce of broad (EWIIM ~ 10000 kins 1) permitted ines in the period July - October 1998: they explained. it as the result of supernova explosions im the eircuummuclear starburst. rather than as a chanee of the reddening iu the orus which would allow to sce the immer nuclear regions.," reported the appearance of broad (FWHM $\sim$ 10000 km $^{-1}$ ) permitted lines in the period July - October 1998: they explained it as the result of supernova explosions in the circumnuclear starburst, rather than as a change of the reddening in the torus which would allow to see the inner nuclear regions."
+ The non-thermal radio ciissiou probably originates in any supernova renimants associated with this starburst. as found bv Forbes Norris (1998)).," The non-thermal radio emission probably originates in many supernova remnants associated with this starburst, as found by Forbes Norris \cite{forbes}) )."
+ Iu the infrared. he 8-13 iau cinission is characterized by the presence 6| polvevelic aromatic lvdrocarbou (PAID) features as revealed by ground-based (Froecl et 1982..," In the infrared, the 8-13 $\mu$ m emission is characterized by the presence of polycyclic aromatic hydrocarbon (PAH) features as revealed by ground-based (Frogel et al. \cite{frogel},"
+ Roche ct al. 198 D)), Roche et al. \cite{roche}) )
+ and. more recently. by ISO SWS observations (Genzel et al. 1998)).," and, more recently, by ISO SWS observations (Genzel et al. \cite{genzel}) )."
+ Their streneth is uuch more typical of starbursts rather than of Ανα in fact. Cenzel et al. (1998))," Their strength is much more typical of starbursts rather than of AGNs: in fact, Genzel et al. \cite{genzel}) )"
+ found that the typical streneth of the qn PAT defined in then Table is 0.01 in ACNs andο3.6 in feature.Maulti-aperturestarbursts: they find a value of 2.5 in NGC ," found that the typical strength of the $\mu$ m PAH feature, defined in their Table 1, is 0.04 in AGNs and 3.6 in starbursts; they find a value of 2.5 in NGC 7582."
+observations performed by Frogcl et al. (1982)), Multi-aperture observations performed by Frogel et al. \cite{frogel}) )
+ showed that the 10 jam enüssiou comes roni a region smaller than 8”. but its origin i« still nuclear.," showed that the 10 $\mu$ m emission comes from a region smaller than $\arcsec$, but its origin is still unclear."
+ Using the FIR to N-rav huuinositv correlation or nonnual aud starburst galaxies. Turner et al. (1997))," Using the FIR to X-ray luminosity correlation for normal and starburst galaxies, Turner et al. \cite{turner}) )"
+ ound that the maxiumui contribution of a starburst to he 0.5-1.5 keV Cluission is 31V, found that the maximum contribution of a starburst to the 0.5-4.5 keV emission is $\sim 34\%$.
+ According to Xue et al. (1998)).," According to Xue et al. \cite{xue}) ),"
+ the soft N-rayv emission (0.5 2 keV) is the sun of scattered emission from the nucleus aux Ες roni the starburst. the latter being ~.," the soft X-ray emission (0.5 - 2 keV) is the sum of scattered emission from the nucleus and emission from the starburst, the latter being $\sim$."
+ Iu their Iubble Space Telescope imaging survey of ucarby ACUNS. Malan et al. (1998))," In their Hubble Space Telescope imaging survey of nearby AGNs, Malkan et al. \cite{malkan}) )"
+ detected dust lanes miming across he nucleus of NGC 7582. at a distance of some huudreds of parsecs: this dust component external to the OLUs iav contribute to obliterate the broad polarized lines whereas he scattered X-ray plotons may still be obxcrved.," detected dust lanes running across the nucleus of NGC 7582, at a distance of some hundreds of parsecs: this dust component external to the torus may contribute to obliterate the broad polarized lines whereas the scattered X-ray photons may still be observed."
+ There is therefore strong evidence tha NGC 7582 contains aburied! ACN vwith the active nucleus residingiu a dusty cuviromment and coexisting with a cireummmclear starburst., There is therefore strong evidence that NGC 7582 contains a `buried' AGN with the active nucleus residing in a dusty environment and coexisting with a circumnuclear starburst.
+ Schachter et al. (1998)), Schachter et al. \cite{schachter}) )
+ suggested that all NLXCs are obscured Ανα., suggested that all NLXGs are obscured AGNs.
+ Because of its prototype nature aud brightness we observed the infrared SED of this object with ISO (Ixessler et al. 1996)), Because of its prototype nature and brightness we observed the infrared SED of this object with ISO (Kessler et al. \cite{kessler}) )
+ to determine the relative contribution of the various components to the tota ΟΠΟΙΟΥ budect., to determine the relative contribution of the various components to the total energy budget.
+ Until recently. the compilation of large statistical samples of clusters of galaxies was a task accomplishable only at optical wavelengths where photographic plates provide both all-sky coverage and sufficient depth to detect clusters at redshifts of 2zz0.8 (te.g.. Abell 1958. Zwicky et 11961-1968. Abell. Corwin Olowin 1989).," Until recently, the compilation of large statistical samples of clusters of galaxies was a task accomplishable only at optical wavelengths where photographic plates provide both all-sky coverage and sufficient depth to detect clusters at redshifts of $z\la 0.3$ (e.g., Abell 1958, Zwicky et 1961–1968, Abell, Corwin Olowin 1989)."
+ Only with the completion of theROSAT AII-Sky Survey (RASS) in 199] (Voges 1992. Trümmper 1993) did unbiased large compilations of X-ray detected clusters become a feasible alternative.," Only with the completion of the All-Sky Survey (RASS) in 1991 (Voges 1992, Trümmper 1993) did unbiased large compilations of X-ray detected clusters become a feasible alternative."
+ To date. three X-ray flux limited cluster samples have been published from RASS data.," To date, three X-ray flux limited cluster samples have been published from RASS data."
+ The all-sky sample of the 242. X-ray Brightest Abell-type Clusters (XBACs) of Ebeling et (1996) was the tirst statistical sample of X-ray bright clusters to emerge from the RASS., The all-sky sample of the 242 X-ray Brightest Abell-type Clusters (XBACs) of Ebeling et (1996) was the first statistical sample of X-ray bright clusters to emerge from the RASS.
+ However. although X-ray flux limited. the XBACs sample is. by design. limited to Abell clusters and thus still affected by the biases inherent in optical cluster surveys.," However, although X-ray flux limited, the XBACs sample is, by design, limited to Abell clusters and thus still affected by the biases inherent in optical cluster surveys."
+ The other two large-scale RASS cluster samples are truly X-ray selected though: the ROSAT Brightest Cluster Sample (BCS. Ebeling et 11998. Paper D comprises 203 X-ray selected clusters in the northern hemisphere. and the RASSI Bright Sample (RASSI-BS. De Grandi et 11999) consists of 130 such clusters in the southern hemisphere.," The other two large-scale RASS cluster samples are truly X-ray selected though: the ROSAT Brightest Cluster Sample (BCS, Ebeling et 1998, Paper I) comprises 203 X-ray selected clusters in the northern hemisphere, and the RASS1 Bright Sample (RASS1-BS, De Grandi et 1999) consists of 130 such clusters in the southern hemisphere."
+ A fourth RASS cluster sample covering most of the southern extragalactic sky is under compilation (Bóhhringer et al.," A fourth RASS cluster sample covering most of the southern extragalactic sky is under compilation (Böhhringer et al,"
+"0.0035(2:0.0006) for v,=90 Km/s. and 0.0012(0.0003) for vo=410 km/s. The ratio of higher energy loss is maximum for lower initial velocities. both in terms of absolute energy and relative energy with respect to E.","$0.0035(\pm 0.0006)$ for $v_\infty =90$ Km/s, and $0.0012(\pm 0.0003)$ for $v_\infty =110$ km/s. The ratio of higher energy loss is maximum for lower initial velocities, both in terms of absolute energy and relative energy with respect to $E_\infty$."
+ The explanation for this is multiple., The explanation for this is multiple.
+ First. for higher velocities the particles go very fast when they reach the bar. and since AE in equation (4)) is proportional to the time the particle spends near the bar. the loss of energy is lower.," First, for higher velocities the particles go very fast when they reach the bar, and since $\Delta E$ in equation \ref{deltaE}) ) is proportional to the time the particle spends near the bar, the loss of energy is lower."
+ Second. the capacity to attract towards the centre particles with high 6 is larger for lower values of the initial velocity.," Second, the capacity to attract towards the centre particles with high $b$ is larger for lower values of the initial velocity."
+ Moreover. Ea; is higher for higher initial velocities. so the rate of energy loss is lower too.," Moreover, $E_\infty$ is higher for higher initial velocities, so the rate of energy loss is lower too."
+" For the cases Vo€30 km/s. it is observed that some particles with b>200 kpe also lose AE<-E,,: therefore. the trapped mass is even higher than the ratio given in eq. (31))."," For the cases $v_\infty \le 30$ km/s, it is observed that some particles with $b>200$ kpc also lose $\Delta E<-E_\infty$; therefore, the trapped mass is even higher than the ratio given in eq. \ref{form200}) )."
+ The dependence is roughly described by the law: The effect of the variation of halo mass is similar to the effect of the variation in the initial velocity because the amount of mass Is proportional to the acceleration. and consequently 1s related to the velocity of the particle.," The dependence is roughly described by the law: The effect of the variation of halo mass is similar to the effect of the variation in the initial velocity because the amount of mass is proportional to the acceleration, and consequently is related to the velocity of the particle."
+ Fig., Fig.
+ 7 illustrates this., \ref{Fig:histo4} illustrates this.
+ A maximum accretion is obtained for the largest mass., A maximum accretion is obtained for the largest mass.
+ Lower masses have lower capacities to trap particles with high 5., Lower masses have lower capacities to trap particles with high $b$ .
+" The values of Roop obtained are: 0.0049(+0.0007) for Miao=0. 0.0092(+£0.0010) for Mpaio=0.2x107M... 0.0220(+0.0015) for My,=O4x107Μ.. 0.0299(+40.0017) for My,=0.6xLO? M... 0.0362(+0.0019) for Mia,=0.8x107M. and 0.0396(+0.0020) for My,=1.0»1077M."," The values of $R_{200}$ obtained are: $\pm $ 0.0007) for $M_{\rm halo}=0$, $\pm $ 0.0010) for $M_{\rm halo}=0.2\times 10^{12}\ {\rm M}_\odot $, $\pm $ 0.0015) for $M_{\rm halo}=0.4\times 10^{12}\ {\rm M}_\odot $, $\pm $ 0.0017) for $M_{\rm halo}=0.6\times 10^{12}\ {\rm M}_\odot $ , $\pm $ 0.0019) for $M_{\rm halo}=0.8\times 10^{12}\ {\rm M}_\odot $, and $\pm $ 0.0020) for $M_{\rm halo}=1.0\times 10^{12}\ {\rm M}_\odot $."
+ It is clear that it is proportional to the halo mass. roughly according to: The gain or loss of mass in the halo due to other mechanisms will not be considered in this paper. but only the mechanism related to the dynamical friction with the bar.," It is clear that it is proportional to the halo mass, roughly according to: The gain or loss of mass in the halo due to other mechanisms will not be considered in this paper, but only the mechanism related to the dynamical friction with the bar."
+ There are other mechanisms of accretion. such as friction with the dise. of particles or clouds that cross it.," There are other mechanisms of accretion, such as friction with the disc of particles or clouds that cross it."
+ There are also ways for the halo to lose mass. such as galaxy stripping in interactions with other galaxies or the IGM. or the escape of particles when they reach a velocity larger than escape velocity due to multiple interactions.," There are also ways for the halo to lose mass, such as galaxy stripping in interactions with other galaxies or the IGM, or the escape of particles when they reach a velocity larger than escape velocity due to multiple interactions."
+ We neglect these other mechanisms because we want to consider only the effect of the bar in increasing the halo mass independently of other effects., We neglect these other mechanisms because we want to consider only the effect of the bar in increasing the halo mass independently of other effects.
+ The reality is more complex. of course. but here we just consider a toy model for rough estimates of the net effect of an increase in the accretion ofmass in barred galaxies with respect to non-barred galaxies.," The reality is more complex, of course, but here we just consider a toy model for rough estimates of the net effect of an increase in the accretion ofmass in barred galaxies with respect to non-barred galaxies."
+Iu addition. our simulations do not include metal lue cooling.,"In addition, our simulations do not include metal line cooling."
+ ? measured the metallicity of a Lyman-lit svstenài and a few super Lymau-limit svstems at x35 ZlobZ.. similar to the average metallicity of damped Lye systems.," \citet{prochter10} measured the metallicity of a Lyman-limit system and a few super Lyman-limit systems at $z\approx3.5$ to be $Z\approx10^{-1.7}~Z_\odot$, similar to the average metallicity of damped $\alpha$ systems."
+ This metallicity is iusuffiieut for metals to be the dominant coolant at Tzr20.000 IK. which requies Z>10!Z. (?).," This metallicity is insufficient for metals to be the dominant coolant at $T\gtrsim20,000~$ K, which requires $Z\gtrsim10^{-1}~Z_\odot$ \citep{wiersma09}."
+ ILlowever. this metallicity sufficient to allow ionized eas to cool to somewhat lower temperatures than in the simulations. where CD;z:20.000 I& at relevant columns.," However, this metallicity sufficient to allow ionized gas to cool to somewhat lower temperatures than in the simulations, where $\langle T\rangle\approx 20,000~$ K at relevant columns."
+ Feedback frou galaxies could result iu the eas within Lyxuiui-linüt svsteuis being shock heated. driving it iuto a collisional rather than photoionization equilibriu.," Feedback from galaxies could result in the gas within Lyman-limit systems being shock heated, driving it into a collisional rather than photoionization equilibrium."
+ If the Lxauau-liuüt svstems were in collisional ionization equilibrimu. their cross section would be inseusitive to I. so that P svould scale with the euussivity.," If the Lyman-limit systems were in collisional ionization equilibrium, their cross section would be insensitive to $\Gamma$, so that $\Gamma$ would scale with the emissivity."
+ Iu our sinulatious with and without ealactic winds. collisional jonizations are unuimuportaut for A~100 eas at τ=6. but become more iniportant with decreasing redshift aud. in fact. set the ionization state for half of A~LOO eas at 2=3.," In our simulations with and without galactic winds, collisional ionizations are unimportant for $\Delta \sim 100$ gas at $z=6$, but become more important with decreasing redshift and, in fact, set the ionization state for half of $\Delta \sim 100$ gas at $z=3$."
+ In fact. we find that collisional ionizatious play sole vole in making our caleulations smaller at οx1 than the predictions based on equation (1)).," In fact, we find that collisional ionizations play some role in making our calculations' $n$ smaller at $z\leq4$ than the predictions based on equation \ref{eqn:relationG}) )."
+ As a final note. if 3 in reality were assteep at NipX GinZ as the observations of ? ο collisional ionization is likely the oulv weaus to siguificautlv weaken the dependence of Toon ε from the scaling im equation (1)).," As a final note, if $\beta$ in reality were assteep at $N_{\rm HI} \lesssim 10^{17.5}~$ $^{-2}$ as the observations of \citet{prochaska10} suggest, collisional ionization is likely the only means to significantly weaken the dependence of $\Gamma$ on $\epsilon$ from the scaling in equation \ref{eqn:relationG}) )."
+ Wile the extragalactic ionizing backeround almost certainly dominates in systems with Nyy< 10120? (?).. local sources of radiation could be nuportaut for higher columumu-deusity svstenis (2)..," While the extragalactic ionizing background almost certainly dominates in systems with $N_{\rm HI}< 10^{17.2}$ $^{-2}$ \citep{miralda05}, local sources of radiation could be important for higher column-density systems \citep{schaye06}."
+ Our calculations do not account for their contribution., Our calculations do not account for their contribution.
+ The analytic model in &L.1 can be generalized to allow for a fraction of tle easat cach density. f(A). to be ionized bv local sources.," The analytic model in \ref{ss:analytics} can be generalized to allow for a fraction of the gasat each density, $f(\Delta)$, to be ionized by local sources."
+ Iu this model. the P. dependence of the emissivitv would oulv be weakened relative to our previous findiugs if PN)|l.ΓΔ) decreased less quickly with A than P(A).," In this model, the $\Gamma$ dependence of the emissivity would only be weakened relative to our previous findings if $P(\Delta)\,[1-f(\Delta)]$ decreased less quickly with $\Delta$ than $P(\Delta)$."
+ ILowever. the opposite secs more likely to hold. i... that local sources prefercutially ionize the denser regious near the sites of star formation.," However, the opposite seems more likely to hold, i.e., that local sources preferentially ionize the denser regions near the sites of star formation."
+ Reiouization is only crudely iuelided im our siauulatiouns. with the intergalactic gas beige reionized homogencously at 2%10 in these calculations.," Reionization is only crudely included in our simulations, with the intergalactic gas being reionized homogeneously at $z\approx10$ in these calculations."
+ However. reiomization could cud as late as :z5.5 (7)..," However, reionization could end as late as $z\approx 5.5$ \citep{mesinger10}."
+ Our results iav be affected if reionizatiou occured at lower redshifts than is assunued im the simulations., Our results may be affected if reionization occurred at lower redshifts than is assumed in the simulations.
+" The relaxation timescale after a system was relonized ds roughly equal to the Jeans length divided by the sound speed. or ΠΠ)tA7,"," The relaxation timescale after a system was reionized is roughly equal to the Jeans length divided by the sound speed, or $H(z)^{-1}\,\Delta^{-1/2}$."
+ Tt takes AzmAUT(gp) after a region was roiouized for the eas to have relaxed.," It takes $\Delta z\approx\Delta^{-1/2}\,(1+z)$ after a region was reionized for the gas to have relaxed."
+ For the A~100 systems of interest. this corresponds to Αι.ο|i).," For the $\Delta \sim 100$ systems of interest, this corresponds to $\Delta z\sim 0.1 \, (1+z)$."
+" After a larec-scale region was reionized. the local P(A) would decrease. starting at the highest densities and evolving to lower densities,"," After a large-scale region was reionized, the local $P(\Delta)$ would decrease, starting at the highest densities and evolving to lower densities."
+" This inside-out process would likely steepeu the scaling between DP. aud time interval of ⋅⋅Az 21/2 A,€;a(L|:) after a region was relonized."," This inside-out process would likely steepen the scaling between $\Gamma$ and $\epsilon$ a time interval of $\Delta z\gtrsim\Delta_i^{-1/2}\,(1+z)$ after a region was reionized."
+ These considerations would also result iu our caleulations underestimating the chuupiuess of the ionized gas just after reionization. aud thereby the Cluissivity. at fixed T.," These considerations would also result in our calculations underestimating the clumpiness of the ionized gas just after reionization, and thereby the emissivity, at fixed $\Gamma$."
+ The abundance of Lyiuau-liuüt svsteius in our fiducial calculation is consistent with observatious at the 20 level," The abundance of Lyman-limit systems in our fiducial calculation is consistent with observations at the $2 \,\sigma$ level."
+" We find that adding feedback frou, galactic winds or lowering DP improves the aerecment.", We find that adding feedback from galactic winds or lowering $\Gamma$ improves the agreement.
+ This conparisou tests whether our cosmological simulations reproduce the properties of eas at the outskirts of galactic halos., This comparison tests whether our cosmological simulations reproduce the properties of gas at the outskirts of galactic halos.
+ However. our simulations may uuderproduce the munber of Ny~101? 7? systems compared with observations.," However, our simulations may underproduce the number of $N_{\rm HI} \sim 10^{19}~$ $^{-2}$ systems compared with observations."
+ If real. this discrepancy could indicate that the gas is cooler and. as a result. deuser in these systenis than in the sinulatious owing to additional cooling. or that the imipact of feedback is more prominent than in the considered wind model.," If real, this discrepancy could indicate that the gas is cooler and, as a result, denser in these systems than in the simulations owing to additional cooling, or that the impact of feedback is more prominent than in the considered wind model."
+ Both observations aud numerical simulations sugeest a steep slope for FCN) with §zmLs at Aym1017cm 7., Both observations and numerical simulations suggest a steep slope for $f(N_{\rm HI})$ with $\beta\approx1.8$ at $N_{\rm HI}\lesssim10^{17.5}~$ $^{-2}$.
+" We showed that this steepness implies a strong relationship between the extragalactic r-ioniziug background aud the ionizing enuüssivitv of the sources,", We showed that this steepness implies a strong relationship between the extragalactic -ionizing background and the ionizing emissivity of the sources.
+ At 53. our calculations find that DP scales as € to the power of 2.2.02.1.," At $z=3$, our calculations find that $\Gamma$ scales as $\epsilon$ to the power of $2.2-2.4$."
+ The near constancy of D that is measured at ;22.L (2?) puts strimecut requirements on the evolution of the total comoving ioniziug Clussivity owing to this strong scale. requiring it to be roughly coustaut.," The near constancy of $\Gamma$ that is measured at $z\approx2-4$ \citep{bolton07, faucher08b} puts stringent requirements on the evolution of the total comoving ionizing emissivity owing to this strong scaling, requiring it to be roughly constant."
+ This is au interesting finding since the ionizing background is believed to transition from being dominated by quasars to being dominated by star-forming galaxies over this redshift iuterval (6.9.. 2)).," This is an interesting finding since the ionizing background is believed to transition from being dominated by quasars to being dominated by star-forming galaxies over this redshift interval (e.g., \citealt{faucher08b}) )."
+ The scaling strengthens with redshift in our calculations. with P proportional to € to the 3.8RE power at +=6.," The scaling strengthens with redshift in our calculations, with $\Gamma$ proportional to $\epsilon$ to the $3.8-4.5$ power at $z=6$."
+ A strong scaling is related to the claim that the IGAL chuuping factor after reionization was snall. independent of A; (?).. 5a.ice both require 523 for the high-density power-law index of the gas deusity distribution.," A strong scaling is related to the claim that the IGM clumping factor after reionization was small, independent of $\Delta_i$ \citep{pawlik09}, since both require $\gamma \approx 3$ for the high-density power-law index of the gas density distribution."
+ Interestingly. ?.— detected a sharp decrease in the Lye forest transmission at τzc6.," Interestingly, \citet{fan06} detected a sharp decrease in the $\alpha$ forest transmission at $z\approx6$."
+ If interpreted as a change in D. 7. estimated that a factor of at least 2 change was required over ZAz0.5.," If interpreted as a change in $\Gamma$, \citet{fan06} estimated that a factor of at least $2$ change was required over $\Delta z\approx0.5$."
+ It does uot seen plausible that the cuussivity of galaxies changed by sucli a large factor over this interval. which constitutes 7% of the ITubble time at 2= 6.," It does not seem plausible that the emissivity of galaxies changed by such a large factor over this interval, which constitutes $7\%$ of the Hubble time at $z=6$ ."
+ Our simulations fine that a more modest z20% chanee in the comoving cuissivity at 2&6 over Atzx0.5 would result iu a factor of 2 change in D., Our simulations find that a more modest $\approx 20$ change in the comoving emissivity at $z\approx6$ over $\Delta z\approx0.5$ would result in a factor of $2$ change in $\Gamma$.
+ Estimates for the variance in D(:) between widely separated regions sugeest that this effect could operate on this timescale (?).., Estimates for the variance in $\Gamma(z)$ between widely separated regions suggest that this effect could operate on this timescale \citep{mesinger09}.
+ The quick evolution in E from the mechauisiu discussed here is nof necessarily associated with the eund of reiouization. defined as the epoch when ueutral patches were present in the low-deusity IGM so that A;~ 1.," The quick evolution in $\Gamma$ from the mechanism discussed here is not necessarily associated with the end of reionization, defined as the epoch when neutral patches were present in the low-density IGM so that $\Delta_i \sim 1$ ."
+ Iu all of our caleulatious. A;z9 1. with A;~50 at :zm6 and for T=10sy + approximately what is weasured at this 2 using the Lya and Ly. forests.," In all of our calculations, $\Delta_i \gg 1$ , with $\Delta_i \sim 50$ at $z\approx 6$ and for $\Gamma = 10^{-13}~$ $^{-1}$ – approximately what is measured at this $z$ using the $\alpha$ and $\beta$ forests."
+ In addition. the mean free path of 1 Ry photous is 8 eMpe for this case.," In addition, the mean free path of $1~$ Ry photons is $8~$ cMpc for this case."
+" This is still larger than the 0.5(1.1) cAIpe average separationsof galaxies residing in >107 (> 107?) 44, halos iu the asstmed cosmology.", This is still larger than the $0.5~(1.1)~$ cMpc average separationsof galaxies residing in $>10^8$ $>10^9$ ) $M_\odot$ halos in the assumed cosmology.
+ Therefore. it does not satisty the simplest criterion for reiouization that the mean free path was shorter than the separation of the," Therefore, it does not satisfy the simplest criterion for reionization that the mean free path was shorter than the separation of the"
+"no more success, if allowing for a differential extinction between the star and the galaxy would reduce the effect and improve the fit.","no more success, if allowing for a differential extinction between the star and the galaxy would reduce the effect and improve the fit."
+" 'The close examination of the fits revealed the probable presence of additional continuum light, in particular on the south side of the galaxy, which is more affected by the star."," The close examination of the fits revealed the probable presence of additional continuum light, in particular on the south side of the galaxy, which is more affected by the star."
+ This pushes the solution to a lower metallicity and degrades the quality of the fit., This pushes the solution to a lower metallicity and degrades the quality of the fit.
+" We then compared the light profiles along the slit for our target and for a template star observed with the same setup, and found that the contamination is actually higher than the one inferred from the spectroscopic decomposition."," We then compared the light profiles along the slit for our target and for a template star observed with the same setup, and found that the contamination is actually higher than the one inferred from the spectroscopic decomposition."
+" This provides additional evidence for pollution by diffuse light, which was not (or not perfectly) dispersed."," This provides additional evidence for pollution by diffuse light, which was not (or not perfectly) dispersed."
+ The scattered light represents of the incident light., The scattered light represents of the incident light.
+" Unfortunately, we did not have adequate observations to check if the PSF before the slit has indeed lower wings (the acquisition images are saturated)."," Unfortunately, we did not have adequate observations to check if the PSF before the slit has indeed lower wings (the acquisition images are saturated)."
+" Based on this suspicion of diffuse light, we fitted models with an additional free additive term."," Based on this suspicion of diffuse light, we fitted models with an additional free additive term."
+" However, such a component is naturally degenerate with the metallicity, and the fit became unstable."," However, such a component is naturally degenerate with the metallicity, and the fit became unstable."
+" It returned on average higher and super-solar [Fe/H], (as expected).", It returned on average higher and super-solar ${\rm [Fe/H]}_g$ (as expected).
+ We therefore added constraints., We therefore added constraints.
+" We modified the model of Equation 1 by replacing the stellar component, S, by a combination (1—a)xS+ax C, where C is a constantspectrum representing the undispersed light and a the corresponding fraction."," We modified the model of Equation \ref{eqn:main} by replacing the stellar component, $S$, by a combination $(1-a) \times S + a \times
+C$ , where $C$ is a constantspectrum representing the undispersed light and $a$ the corresponding fraction."
+" We adjusted a as the function of the distance to the star along the slit to match the retrieved spectroscopic contamination, f, with the photometric contamination obtained comparing our observation with one of a template star."," We adjusted $a$ as the function of the distance to the star along the slit to match the retrieved spectroscopic contamination, $f$, with the photometric contamination obtained comparing our observation with one of a template star."
+ We tried if using a different spectral energy distribution in place of the constant C affected the result and the quality of the fit., We tried if using a different spectral energy distribution in place of the constant $C$ affected the result and the quality of the fit.
+ We considered either the same distribution as the star (using a smooth spectrum of it) or a spectrum biased to the blue (because the undispersed light may contaminate more the blue part of the spectrum)., We considered either the same distribution as the star (using a smooth spectrum of it) or a spectrum biased to the blue (because the undispersed light may contaminate more the blue part of the spectrum).
+ But we found no significant difference between these choices and adopted the additive constant., But we found no significant difference between these choices and adopted the additive constant.
+" Because the age was apparently homogeneous, we fixed it to 12 Gyr to reduce the degree of freedom."," Because the age was apparently homogeneous, we fixed it to 12 Gyr to reduce the degree of freedom."
+ 'The final profiles of the kinematics and of the stellar population are presented in Fig. 2.., The final profiles of the kinematics and of the stellar population are presented in Fig. \ref{fig:profile}.
+" We rejected the points in the range -13<z«19 ""(15«r42 ""on the southern side), where the contamination from the foreground star exceeds the galactic light by more than a factor 2."," We rejected the points in the range $-1.3 < x < 1.9$ $1.5 < r < 4.2$ on the southern side), where the contamination from the foreground star exceeds the galactic light by more than a factor 2."
+ An example of the decomposition is shown on Fig., An example of the decomposition is shown on Fig.
+" 3 for the spectrum at the location z=3,r5.6"",, where the two contributions are comparable."," \ref{fig:fit} for the spectrum at the location $x=3, 
+r=5.6$, where the two contributions are comparable."
+ The residuals are consistent with the noise computed from the photon statistics., The residuals are consistent with the noise computed from the photon statistics.
+" 'The kinematical, and now the metallicity, profiles are symmetric."," The kinematical, and now the metallicity, profiles are symmetric."
+ The systemic velocity is 7574:6kms-!., The systemic velocity is $757\pm6$.
+". The peak of rotation velocity, πας=107+ 3, is reached at 2 along the major axis (i. e. at a galactocentric"," The peak of rotation velocity, $
+V_{\rm max} = 107\pm3$ , is reached at 2 along the major axis (i. e. at a galactocentric"
+are capable of supporting stellar svstems with realistic kinematics.,are capable of supporting stellar systems with realistic kinematics.
+ ο demonstrated that all of the classical ATW satellites for which high quality kinematic data are available are consistent with having formed in dark matter subhalos selected from the high resolution N-body simulations of the Aquarius project. (2)..., \citet{StrigariFrenkWhite2010} demonstrated that all of the classical MW satellites for which high quality kinematic data are available are consistent with having formed in dark matter subhalos selected from the high resolution N-body simulations of the Aquarius project \citep{Springel2008}.
+ However. successful models must satisfy both constraints. producing realistic Luminosity l'unctions forming stars in potentials like those inferred [rom observations.," However, successful models must satisfy both constraints, producing realistic luminosity functions forming stars in potentials like those inferred from observations."
+ ? have recently argued that the most massive subhalos in high resolution simulations of cold clark matter halos are too concentrated to be able to host. the brightest observed satellites of the Milkv Way., \citet{Boylan-Kolchin2011} have recently argued that the most massive subhalos in high resolution simulations of cold dark matter halos are too concentrated to be able to host the brightest observed satellites of the Milky Way.
+ ο have shown that subhalos formed. from warm. rather than cold. dark matter have suitably low concentrations. but both they and DBovlan-Ixolchin ct al.," \citet{Lovell2011} have shown that subhalos formed from warm, rather than cold, dark matter have suitably low concentrations, but both they and Boylan-Kolchin et al."
+ emphasise that other solutions to the cliscrepaney are possible., emphasise that other solutions to the discrepancy are possible.
+ One promising possibility is the mechanism originally proposed by Navarro. Eke Frenk (199G6a). whereby the rapid. expulsion of dense central gas in a starburst can unbind the inner parts of the halo. significantly reducing its concentration.," One promising possibility is the mechanism originally proposed by Navarro, Eke Frenk (1996a), whereby the rapid expulsion of dense central gas in a starburst can unbind the inner parts of the halo, significantly reducing its concentration."
+ We find an example of this process in one of the subhalos formed in the simulations analvsed in this paper., We find an example of this process in one of the subhalos formed in the simulations analysed in this paper.
+ From a theoretical perspective.ilio hvdrodynamic simulations are uniquely well-suited. to investigating the cllects of galaxy formation on the dark. matter halos of satellite galaxies.," From a theoretical perspective, hydrodynamic simulations are uniquely well-suited to investigating the effects of galaxy formation on the dark matter halos of satellite galaxies."
+ It has been. known for some time that barvons may significantl alter the behaviour of dark matter on some scales., It has been known for some time that baryons may significantly alter the behaviour of dark matter on some scales.
+ Dissipative processes such as gas cooling. star formation and feedback decouple the dynamical evolution of the barvons from that of the dark matter.," Dissipative processes such as gas cooling, star formation and feedback decouple the dynamical evolution of the baryons from that of the dark matter."
+ The resulting change in the shape of the overall potential in turn allects the phase space structure of the dark matter halo., The resulting change in the shape of the overall potential in turn affects the phase space structure of the dark matter halo.
+ Central. concentrations of cold. barvonic matter can induce an adiabatic. radial contraction of the central regions (?7).. while the opposite οσοι can he achieved if dense clumps of barvonic material heat the central distribution of dark matter (227).. or i£ the blowout mechanism of Navarro. Eke and Frenk is effective (seealso.22????)..," Central concentrations of cold baryonic matter can induce an adiabatic, radial contraction of the central regions \citep{Blumenthal1986,Gnedin2004}, while the opposite effect can be achieved if dense clumps of baryonic material heat the central distribution of dark matter \citep{MoMao2004,Mashchenko2006,Mashchenko2008}, or if the blowout mechanism of Navarro, Eke and Frenk is effective \citep[see
+also][]{GelatoSommerlarsen1999,GnedinZhao2002,MoMao2004,ReadGilmore2005,Governato2010,PontzenGovernato2011}."
+ This latter mechanism was originally proposed. as à means to erase the central dark matter cusps in. dwarf. galaxies. though whether or not such ‘cored’ profiles are required. by the observations remains à matter of ongoing debate (?????)..," This latter mechanism was originally proposed as a means to erase the central dark matter cusps in dwarf galaxies, though whether or not such `cored' profiles are required by the observations remains a matter of ongoing debate \citep{Goerdt2006,SanchezSalcedo2006,Strigari2006,Gilmore2007,Walker2009}."
+ If barvons really co modify. the dark matter in satellites on sub-kpe scales significantly. then the value of studying chwarl galaxies with post-processed N-body simulations may να very limited.," If baryons really do modify the dark matter in satellites on sub-kpc scales significantly, then the value of studying dwarf galaxies with post-processed N-body simulations may be very limited."
+ In this paper. we make use of a model that has already had success in reproducing some properties of he Local Croup satellites. including the shape anc approximate normalisation of their Luminosity function ane he metallicity-luminosity relation (7)..," In this paper, we make use of a model that has already had success in reproducing some properties of the Local Group satellites, including the shape and approximate normalisation of their luminosity function and the metallicity-luminosity relation \citep{Okamoto2010a}."
+ In we outline the details of our simulations. including the initia conditions. simulation code ancl modelling of various kev owvonie physical processes.," In we outline the details of our simulations, including the initial conditions, simulation code and modelling of various key baryonic physical processes."
+ In.3.. we expand on he theoretical predictions of the model by examining wha οσοι barvons have had on the dark matter profiles of satellites.," In, we expand on the theoretical predictions of the model by examining what effect baryons have had on the dark matter profiles of satellites."
+ We perform. tests in to ensure tha kev properties of our satellite population do not depend on he numerical resolution., We perform tests in to ensure that key properties of our satellite population do not depend on the numerical resolution.
+ In Dowe CODipare severa observable and derived properties of the simulated satellites o Local Group data., In we compare several observable and derived properties of the simulated satellites to Local Group data.
+ Finally. in we discuss the evolution of one particularly interesting satellite in the simulation. which is dominated by its stellar component at >=0.," Finally, in we discuss the evolution of one particularly interesting satellite in the simulation, which is dominated by its stellar component at $z=0$."
+ Our main results are sunimerisecd in7., Our main results are summarised in.
+". To investigate the properties of a simulated: MW-satellite system. we select one of the six halos from the Aquarius project described in 2.. halo ""€"" in their labelling system."," To investigate the properties of a simulated MW-satellite system, we select one of the six halos from the Aquarius project described in \citet{Springel2008}, halo `C' in their labelling system."
+ These halos were extracted from a cosmological simulation in a cube of comoving volume (100Mpc)? and were chosen to have masses close to that of the Milky Way (~1012AZ.) and avoid dense environments (no neighbour exceeding half its. mass within.. 1.1 “Alpe) (?).., These halos were extracted from a cosmological simulation in a cube of comoving volume $\rm (100Mpc)^{3}$ and were chosen to have masses close to that of the Milky Way $\sim10^{12}M_{\odot}$ ) and avoid dense environments (no neighbour exceeding half its mass within $1h^{-1}$ Mpc) \citep{Navarro2010}.
+: As dn Aquarius. we emplov a ‘zoom resimulation technique. with higher mass boundary particles used. to model the Large scale potential ancl lower mass particles in à ~5hi Alpe region surrouncling the target halo.," As in Aquarius, we employ a `zoom' resimulation technique, with higher mass boundary particles used to model the large scale potential and lower mass particles in a $\sim5 h^{-1}$ Mpc region surrounding the target halo."
+ Extra power is ασσα to the initial particle distribution on small scales in the high resolution region. as described hy 7..," Extra power is added to the initial particle distribution on small scales in the high resolution region, as described by \citet{Frenk1996}."
+" We assume a cosmology. with parameters (2,=0.25. Oy=0.75. ο,= 0.045. 08—0.9. n;= Land 4o=100hkms!Mpe+T3kms+Alpe +."," We assume a cosmology, with parameters $\Omega_{m}=0.25$, $\Omega_{\Lambda}=0.75$, $\Omega_{b}=0.045$ , $\sigma_{8}=0.9$, $n_{s}=1$ and $H_{0}=100h {\rm kms}^{-1}{\rm Mpc}^{-1}=73{\rm
+ kms}^{-1}{\rm Mpc}^{-1}$ ."
+ The highest resolution realisation of halo C in Aquarius had a dark matter particle mass of 1.410M., The highest resolution realisation of halo C in Aquarius had a dark matter particle mass of $\rm 1.4\times10^4M_{\odot}$.
+ However. the extra computational time associated with hyvdrodynamic simulations makes such a resolution impractical: our highest resolution instead. corresponds to a dark matter particle mass of ~2.6.107M. and an initial gas particle nass of 5.8«.107M...," However, the extra computational time associated with hydrodynamic simulations makes such a resolution impractical; our highest resolution instead corresponds to a dark matter particle mass of $\rm
+\sim2.6\times10^5M_{\odot}$ and an initial gas particle mass of $\rm
+5.8\times10^4M_{\odot}$."
+ In order to conduct. convergence studies. we also simulated. the halo at. two lower resolutions. with particle masses S anc ~64 times larger.," In order to conduct convergence studies, we also simulated the halo at two lower resolutions, with particle masses $\sim8$ and $\sim64$ times larger."
+ We adopt the same naming convention as ?.. labelling the three runs (in order of decreasing resolution) Aq-C-4. Aq-C-5 and Aq-C-6.," We adopt the same naming convention as \citet{Springel2008}, labelling the three runs (in order of decreasing resolution) Aq-C-4, Aq-C-5 and Aq-C-6."
+ lists the numerical parameters of cach simulation., lists the numerical parameters of each simulation.
+ Our simulation code is based on an early version of the PAI-Pree-SPLL codeGADGEI-3., Our simulation code is based on an early version of the PM-Tree-SPH code.
+. Barvonic processes are mocelled as described in 7.. with a number of modifications designed. to improve the treatment of supernovac-driven winds.," Baryonic processes are modelled as described in \citet{Okamoto2010a}, with a number of modifications designed to improve the treatment of supernovae-driven winds."
+ In the following subsections. we summarise some of the most important. features of the code with emphasis on aspects that have the greatest impact on satellite formation.," In the following subsections, we summarise some of the most important features of the code with emphasis on aspects that have the greatest impact on satellite formation."
+ Racliative processes in our model are implemented. as described in ο and include inverse Compton scattering of CALIB photons. thermalBremsstrahlung. atomic line cooling and. photoionisation heating from. IIvdrogen ancl Helium.," Radiative processes in our model are implemented as described in \citet{WiersmaSchayeSmith2009} and include inverse Compton scattering of CMB photons, thermalBremsstrahlung, atomic line cooling and photoionisation heating from Hydrogen and Helium."
+ All eas in the simulation volume is ionised and heated, All gas in the simulation volume is ionised and heated
+"The two fine structure lines of aat 158um and aat um, are the strongest cooling lines of the ISM, carrying up to a few percent of the total energy emitted from galaxies in the far-infrared wavelengths.","The two fine structure lines of at $\,\mu$ m and at $\mu$ m, are the strongest cooling lines of the ISM, carrying up to a few percent of the total energy emitted from galaxies in the far-infrared wavelengths."
+ The lline lies KK above the ground state with a critical density for collisions with H of 3x10? ccm? (Kaufmanetal.1999)., The line lies K above the ground state with a critical density for collisions with H of $3\times10^3$ $^{-3}$ \citep{Kaufman1999}.
+". While both lines are thought to trace photon-dominated regions (PDRs) at the FUV irradiated surfaces of molecular clouds, it has been realized early-on that a non-negligible fraction of the eemission may stem from the ionized medium (Heiles1994)."," While both lines are thought to trace photon-dominated regions (PDRs) at the FUV irradiated surfaces of molecular clouds, it has been realized early-on that a non-negligible fraction of the emission may stem from the ionized medium \citep{Heiles1994}."
+". Owing to its higher upper energy level KK) and higher critical density of ~5x10°cmcub,, the um lline is a more dominant coolant in warmer and denser neutral regions (e.g.Rólligetal.2006)."," Owing to its higher upper energy level K) and higher critical density of $\simeq 5\times
+10^5$, the $\,\mu$ line is a more dominant coolant in warmer and denser neutral regions \citep[e.g.][]{Roellig2006}."
+". The lline together with the lline are diagnostics to infer the physical conditions in the gas, its temperatures, densities, and radiation fields, by comparing the intensities and their ratios with predictions of PDR models (e.g.,Tielens&Hollenbach1985;Wolfireetal.1990;Kaufmanetal.1999;Róllig2007;Ferland 1998)."," The line together with the line are diagnostics to infer the physical conditions in the gas, its temperatures, densities, and radiation fields, by comparing the intensities and their ratios with predictions of PDR models \citep[e.g.,][]{TielensHollenbach1985,
+Wolfire1990, Kaufman1999, Roellig2007, Ferland1998}."
+. Previous observational studies of eemission from external galaxies include the statistical studies by Crawfordetal.(1985);Stacey(1991);," Previous observational studies of emission from external galaxies include the statistical studies by \citet{Crawford1985,stacey1991,malhotra2001}."
+" More recently the eemission from a few individual galaxies e.g., LMC (Israeletal.1996), M51 (Nikolaetal.2001;Kramer2005),, NGC 6946 (Maddenetal.1993;Contursi2002),, M83 (Krameretal.2005),, NGC 1313 (Contursietal.2002),, M 31 (Rodriguez-Fernandezetal.2006),, NGC1097 (Contursietal.2002;Beiráoetal.2010) have been observed."," More recently the emission from a few individual galaxies e.g., LMC \citep{israel1996}, , M51 \citep{nikola2001,kramer2005}, NGC 6946 \citep{Madden1993,Contursi2002}, M83 \citep{kramer2005}, NGC 1313 \citep{Contursi2002}, M 31 \citep{rodriguez2006}, NGC1097 \citep{Contursi2002,Beirao2010} have been observed."
+ These papers explore the origin of the eemission., These papers explore the origin of the emission.
+" Of these studies, the study of LMC by (1996) was at a resolution of 16 pc and Rodriguez-Fernandez resolved the spiral arms of M 31 at spatial scales of 300 pc."," Of these studies, the study of LMC by \citet{israel1996} was at a resolution of 16 pc and \citet{rodriguez2006} resolved the spiral arms of M 31 at spatial scales of 300 pc."
+" In addition to aand663um,, there are several additional FIR lines, the mid-/ CO transitions, and lines of single and double ionized N and O which provide information about the gas density (multiple transitions of the same ions), hardness of the stellar radiation field (ratio of intensities of two ionization states of the same species), and ionizing flux."," In addition to and, there are several additional FIR lines, the $J$ CO transitions, and lines of single and double ionized N and O which provide information about the gas density (multiple transitions of the same ions), hardness of the stellar radiation field (ratio of intensities of two ionization states of the same species), and ionizing flux."
+" In particular the llines at 122 and um aallow us to estimate the density of the ionized gas, which is a key parameter to model the eemission stemming from the ionized gas."," In particular the lines at 122 and $\,\mu$ allow us to estimate the density of the ionized gas, which is a key parameter to model the emission stemming from the ionized gas."
+" Although much detailed information can be obtained by studying nearby (Milky Way) sites of star formation, a more comprehensive view is possible with a nearby moderately inclined galaxy such as M33."," Although much detailed information can be obtained by studying nearby (Milky Way) sites of star formation, a more comprehensive view is possible with a nearby moderately inclined galaxy such as M33."
+" 333 is a nearby, moderately metal poor late-type spiral galaxy with no bulge or ring, classified as SA(s)cd."," 33 is a nearby, moderately metal poor late-type spiral galaxy with no bulge or ring, classified as SA(s)cd."
+ It is the 3rd largest member of the Local Group., It is the 3rd largest member of the Local Group.
+" Its mass, size, and average metallicity are similar to those of the Large Magellanic Cloud (LMC)."," Its mass, size, and average metallicity are similar to those of the Large Magellanic Cloud (LMC)."
+ M33 hosts some of the brightest ccomplexes in the Local Group., M33 hosts some of the brightest complexes in the Local Group.
+ 6604 is the second brightest rregion after DDoradus in the LMC., 604 is the second brightest region after Doradus in the LMC.
+ Its inclination Gi= 56°)(Regan&Vogel1994) yields a small line-of-sight, Its inclination $i=56^\circ$ )\citep{Regan1994} yields a small line-of-sight
+"well bevond the initial /,,,,. a substantial fraction of the core would be eroded and mixed with the envelope.","well beyond the initial $R_{core}$, a substantial fraction of the core would be eroded and mixed with the envelope."
+" However. if after the impact. 77,=Deoreo-ATore<Deap ΟΥ the region with T'>T5, is confined to small radius. the core structure would be mostly preserved even though convection may be able to penetrate through (he initial core-envelope interlace."," However, if after the impact, $T_{\rm core}
+^\prime = T_{\rm core} + \Delta T_{\rm core} < T_{\rm evap}$ or the region with $T> T_{2 p}$ is confined to small radius, the core structure would be mostly preserved even though convection may be able to penetrate through the initial core-envelope interface."
+" We now consider (he possibility of GIMs onto gas giants with cool (Zu,<Leona) cores."," We now consider the possibility of GIMs onto gas giants with cool $T_{\rm
+core} < T_{\rm cond}$ ) cores."
+ With sufficiently massive and energetic impactor. head-on collisions lead to Ty).>τομ.," With sufficiently massive and energetic impactor, head-on collisions lead to $T_{\rm core} ^\prime > T_{\rm evap}$."
+ The subsequent evolution is similar to that discussed above. albeit a modest faction of internal enerev must be used to compensate for the latent heat curing the melting and evaporation of the core.," The subsequent evolution is similar to that discussed above, albeit a modest fraction of internal energy must be used to compensate for the latent heat during the melting and evaporation of the core."
+ Less energetic impacts may nol be able to evaporate the core. in which case. ihe temperature distribution of the core would be determined bx (he energy equation (3)) including heat transport by conduction.," Less energetic impacts may not be able to evaporate the core, in which case, the temperature distribution of the core would be determined by the energy equation \ref{eqn:energy}) ) including heat transport by conduction."
+" Provided the new LT7Ti,423 the core-envelope interlace. the molten heavy elements would be entrained into the convective eddies on the erowth time scale for Ixelvin-Ilelmholtz instability."," Provided the new $T > T_{\rm cond}$ at the core-envelope interface, the molten heavy elements would be entrained into the convective eddies on the growth time scale for Kelvin-Helmholtz instability."
+ This rate may be limited by the poorly determined material properties such as surface tension and phase separation., This rate may be limited by the poorly determined material properties such as surface tension and phase separation.
+" However. if T>Ts, at the core-envelope interface. mixing between the heavy elements and hydrogen mav be much more efficient."," However, if $T > T_{2
+p}$ at the core-envelope interface, mixing between the heavy elements and hydrogen may be much more efficient."
+ Since convection in the envelope extents to the interlace. the mixed gas is redistributed throughout the convective region with T>Ts).," Since convection in the envelope extents to the interface, the mixed gas is redistributed throughout the convective region with $T > T_{2 p}$."
+ The main objective of this paper is to suggest that eiat impacts and merger process may have led to different core-envelope structure between Jupiter and Saturn aud (he super-solar metallicitv in their gaseous envelopes., The main objective of this paper is to suggest that giant impacts and merger process may have led to different core-envelope structure between Jupiter and Saturn and the super-solar metallicity in their gaseous envelopes.
+ We have highlighted several avenues of GIMs by. embryos onto gas giants., We have highlighted several avenues of GIMs by embryos onto gas giants.
+ The supply of these Luge solid buikling blocks include:situ embryos during the final phase of gas eijant formation. and terrestrial bodies cleared by sweeping secular resonances and ανασα] instabiliües.," The supply of these large solid building blocks include: embryos during the final phase of gas giant formation, and terrestrial bodies cleared by sweeping secular resonances and dynamical instabilities."
+ There are also several processes which can lead (o the merger of eas giants. including orbit crossing triggered bv run-away migration. and long-term dynamical instability.," There are also several processes which can lead to the merger of gas giants, including orbit crossing triggered by run-away migration, and long-term dynamical instability."
+ At their present-day orbits. the velocity dispersion of anv residual embryos or planetesimals is less (han the eas giants! surface escape speed.," At their present-day orbits, the velocity dispersion of any residual embryos or planetesimals is less than the gas giants' surface escape speed."
+ The GIMSs between any residual planetesimnals. embryos. or hvpothetic gas giants and long-period eas giants are likely to be parabolic and abtain limited amount of energy.," The GIMs between any residual planetesimals, embryos, or hypothetic gas giants and long-period gas giants are likely to be parabolic and attain limited amount of energy."
+ This modest velocity dispersion also enlarges eas giants eravitational cross section., This modest velocity dispersion also enlarges gas giants' gravitational cross section.
+ The discussions above indicate (hat embryos wilh a mass of several Af can penetrate through the envelope and reach (he core of gas giants through. head-on collisions., The discussions above indicate that embryos with a mass of several $M_\oplus$ can penetrate through the envelope and reach the core of gas giants through head-on collisions.
+ Η these, If these
+"possible contribution of an ""external"" density distribution (e.g.. associated with a co-existing dark matter halo). is a classical problem of stellar dynamics. with important analogues in the study of collisionless plasmas.","possible contribution of an “external"" density distribution (e.g., associated with a co-existing dark matter halo), is a classical problem of stellar dynamics, with important analogues in the study of collisionless plasmas."
+ In this paper we follow the approach based on the Jeans theorem to deal with the halo matter. while the disk of visible matter 1s taken to provide an external field.," In this paper we follow the approach based on the Jeans theorem to deal with the halo matter, while the disk of visible matter is taken to provide an external field."
+ In this approach. the starting point is the identification. of an appropriate form for the one-particle distribution function in phase space: in our case this is the isothermal assumption for the dark halo.," In this approach, the starting point is the identification of an appropriate form for the one-particle distribution function in phase space; in our case this is the isothermal assumption for the dark halo."
+ The adopted distribution function then generates a @-dependent density in physical space. thus leading. inserted as a source in the Poisson equation. to a nonlinear problem in Φ.," The adopted distribution function then generates a $\Phi$ -dependent density in physical space, thus leading, inserted as a source in the Poisson equation, to a nonlinear problem in $\Phi$."
+ Here @ represents the gravitational potential., Here $\Phi$ represents the gravitational potential.
+ In the past. this method has been used frequently and has led to important insights into the dynamies of galaxies and other stellar systems.," In the past, this method has been used frequently and has led to important insights into the dynamics of galaxies and other stellar systems."
+ Yet the study of self-consistent equilibria in the presence of external fields that break the natural symmetry associated with the one-component problem is only rarely considered (see Bertin Να 2008)., Yet the study of self-consistent equilibria in the presence of external fields that break the natural symmetry associated with the one-component problem is only rarely considered (see Bertin Varri 2008).
+ Therefore. besides the natural application to the discussion of the structure of dark halos around galaxies. this paper has some interest of its own. as a study in stellar dynamics.," Therefore, besides the natural application to the discussion of the structure of dark halos around galaxies, this paper has some interest of its own, as a study in stellar dynamics."
+ The structure of the paper is the following., The structure of the paper is the following.
+ In Sect., In Sect.
+ 2 we comment on the so-called disk-halo conspiracy., 2 we comment on the so-called disk-halo conspiracy.
+ Here we identify an empirical correlation which brings out a new and stronger conspiracy than the one associated with the flatness of rotation curves., Here we identify an empirical correlation which brings out a new and stronger conspiracy than the one associated with the flatness of rotation curves.
+ In Sect., In Sect.
+ 3 we formulate the self-consistent problem of constructing an isothermal spheroid in the presence of a zero-thickness disk., 3 we formulate the self-consistent problem of constructing an isothermal spheroid in the presence of a zero-thickness disk.
+ The actual procedure to calculate the self-consistent models is reported in Appendix A. The properties of self-consistent isothermal halos are then described in Sect., The actual procedure to calculate the self-consistent models is reported in Appendix A. The properties of self-consistent isothermal halos are then described in Sect.
+ 4., 4.
+ Then in Sect., Then in Sect.
+ 5. by modeling first a fiducial rotation curve and later by modeling the specific case of NGC 3198. we show that self-consisteney actually removes the degeneracy that characterizes standard parametric disk-halo decompositions (with additional comments given separately in Appendix B).," 5, by modeling first a fiducial rotation curve and later by modeling the specific case of NGC 3198, we show that self-consistency actually removes the degeneracy that characterizes standard parametric disk-halo decompositions (with additional comments given separately in Appendix B)."
+ Finally. Sect.," Finally, Sect."
+ 6 summarizes the main results of the paper., 6 summarizes the main results of the paper.
+ This article ts all based on the mathematically simple case of zero-thickness disks., This article is all based on the mathematically simple case of zero-thickness disks.
+ A follow-up paper will address the corresponding study of disks with finite thickness., A follow-up paper will address the corresponding study of disks with finite thickness.
+ The rotation curve V(R) of many spiral galaxies is featureless. regular. and smooth: after a central region of approximately linear growth. the rotation velocity reaches a constant value V. and remains flat out to large radii. often well beyond the bright optical disk. [," The rotation curve $V(R)$ of many spiral galaxies is featureless, regular, and smooth: after a central region of approximately linear growth, the rotation velocity reaches a constant value $V_{\infty}$ and remains flat out to large radii, often well beyond the bright optical disk. ["
+But. of course. this is meant to be only a zero-th order description. because even flat rotaton curves are not entirely featureless (Saneist van Albada 1957).,"But, of course, this is meant to be only a zero-th order description, because even flat rotaton curves are not entirely featureless (Sancisi van Albada 1987).]"
+ Since the underlying gravitational field results from the sum of the contributions of different mass components of the system (the stellar disk. the bulge. the gaseous disk. and the dark matter halo). which generally dominate in different regions. the smoothness and lack of features of the rotation curve Is unexpected and poses a physical problem. which ts commonly referred to as theconspiracy.," Since the underlying gravitational field results from the sum of the contributions of different mass components of the system (the stellar disk, the bulge, the gaseous disk, and the dark matter halo), which generally dominate in different regions, the smoothness and lack of features of the rotation curve is unexpected and poses a physical problem, which is commonly referred to as the."
+ When such conspiracy was noted (see van Albada Sancist 1986) most of the emphasis was placed on the fatness of the rotation curves., When such conspiracy was noted (see van Albada Sancisi 1986) most of the emphasis was placed on the flatness of the rotation curves.
+ In reality. 1t has long been noted that standard parametric disk-halo decompositions. based on the superposition of the fields produced by the visible matter and by ahalo.. are subject to a high level of degeneracy.," In reality, it has long been noted that standard parametric disk-halo decompositions, based on the superposition of the fields produced by the visible matter and by a, are subject to a high level of degeneracy."
+ Equally viable fits to the rotation curves are obtained with very different models (1.e. very different values of the mass-to-light ratio of the stellar disk. assumed to be constant within the galaxy). ranging from some characterized by a very light disk and a dominant halo all the way in. to others. close to the so-called maximum disk solution. where the support to the rotation curve inside the optical disk 1s basically due to the visible matter alone. while the halo contribution becomes important ad eventually dominates only outside.," Equally viable fits to the rotation curves are obtained with very different models (i.e. very different values of the mass-to-light ratio of the stellar disk, assumed to be constant within the galaxy), ranging from some characterized by a very light disk and a dominant halo all the way in, to others, close to the so-called maximum disk solution, where the support to the rotation curve inside the optical disk is basically due to the visible matter alone, while the halo contribution becomes important and eventually dominates only outside."
+ In à sense. the existence of such degeneracy actually undermines. or at least softens. the problem of conspiracy apparently raised by the flatness of rotation curves.," In a sense, the existence of such degeneracy actually undermines, or at least softens, the problem of conspiracy apparently raised by the flatness of rotation curves."
+ Of course. there are several constraints that can be used to reduce the level of degeneracy present.," Of course, there are several constraints that can be used to reduce the level of degeneracy present."
+ In addition to those posed by stellar population analyses. we may recall the constraints posed by dynamical stability in the context of spiral structure (Ostriker Peebles 1973; Bertin Lin 1996; see also Widrow et al.," In addition to those posed by stellar population analyses, we may recall the constraints posed by dynamical stability in the context of spiral structure (Ostriker Peebles 1973; Bertin Lin 1996; see also Widrow et al."
+ 2008; and references therein) or warp dynamics (e.g.. see Bertin Mark 1980).," 2008; and references therein) or warp dynamics (e.g., see Bertin Mark 1980)."
+ To give a concrete example of the procedure adopted in standard parametric disk-halo decompositions. we may recall the case of a purely exponential disk with the adoption of the model of a classical pseudo-isothermal for the halo contribution.," To give a concrete example of the procedure adopted in standard parametric disk-halo decompositions, we may recall the case of a purely exponential disk with the adoption of the model of a classical pseudo-isothermal for the halo contribution."
+ In this case the rotation curve is decomposed into Shot=VaTV5., In this case the rotation curve is decomposed into $V_{mod}^2=V_{DM}^2+V_D^2$.
+ Here where €Ες=R//r is the(5) galactocentric-(5) (5)cylindrical radius in units of the exponential scalelength and£ is a dimensionless parameter (see Eq. (16))), Here where $\xi=R/h$ is the galactocentric cylindrical radius in units of the exponential scalelength and $\beta$ is a dimensionless parameter (see Eq. \ref{beta}) ))
+ which measures the weight of the stellar disk., which measures the weight of the stellar disk.
+ In addition. Zo. Ko. 7/4. and Κι denote the standard modified Bessel functions.," In addition, $I_0$, $K_0$, $I_1$, and $K_1$ denote the standard modified Bessel functions."
+ The halo contribution is described by where « is a dimensionless measure of the central dark matter density (see Eq. (19)))., The halo contribution is described by where $\alpha$ is a dimensionless measure of the central dark matter density (see Eq. \ref{alpha}) )).
+ In this example. the degeneracy occurs in the (o.B) parameter space (see Sect. ??)).," In this example, the degeneracy occurs in the $(\alpha, \beta)$ parameter space (see Sect. \ref{decomp}) )."
+ In the following subsection we will show the existence of a much stronger reason for invoking a problem of conspiracy. associated with the fine tuning between the length scale that characterizes the inner growth of the rotation curve and the natural scale length of the visible disk (in the absence of significant contributions from a bulge or from gas. this would be the exponential length).," In the following subsection we will show the existence of a much stronger reason for invoking a problem of conspiracy, associated with the fine tuning between the length scale that characterizes the inner growth of the rotation curve and the natural scale length of the visible disk (in the absence of significant contributions from a bulge or from gas, this would be the exponential length)."
+ This tuning is a direct indication of a subtle interplay between the density distribution of dark and visible matter., This tuning is a direct indication of a subtle interplay between the density distribution of dark and visible matter.
+ If the two length scales were not sufficiently close to the observed correlation. the degeneracy," If the two length scales were not sufficiently close to the observed correlation, the degeneracy"
+redshifted/NW polar lobe of the LIT is essentially invisible.,redshifted/NW polar lobe of the LH is essentially invisible.
+ This implies that the NIZ2 slit. position encounters more extinction in the equator than at other positions. blocking he light from the far side of the nebula.," This implies that the NE2 slit position encounters more extinction in the equator than at other positions, blocking the light from the far side of the nebula."
+" Indeed. this region ocated 2"". 3"" north of the star is relatively dark in optical images and has a high optical depth of cool dust. (Smith et 22003b)."," Indeed, this region located $\arcsec$ $\arcsec$ north of the star is relatively dark in optical images and has a high optical depth of cool dust (Smith et 2003b)."
+ The only sign of the NW lobe is a faint Moteh at O km J|... 27. whieh coincides with a feature in the equatorial skirt at visual wavelengths.," The only sign of the NW lobe is a faint blotch at 0 km $^{-1}$, $\arcsec$, which coincides with a feature in the equatorial skirt at visual wavelengths."
+ This signifies hat bright portions of the skirt may be holes where light can penetrate., This signifies that bright portions of the skirt may be holes where light can penetrate.
+ “Phe SE polar lobe is bright. ancl has a morphology consistent with a slice through a flattened polar rubble.26):," The SE polar lobe is bright, and has a morphology consistent with a slice through a flattened polar bubble.:"
+ Phe kinematic structure at this position already οους from NE2., The kinematic structure at this position already differs from NE2.
+ Part of the recshifted NW. lobe can be seen. although it still sullers more extinction than the blueshifted lobe.," Part of the redshifted NW lobe can be seen, although it still suffers more extinction than the blueshifted lobe."
+ Here the SEE lobe has a more angular or trapezoidal shape than at NIZ2. with straight side walls. pointed corners. and a flat polar cap.," Here the SE lobe has a more angular or trapezoidal shape than at NE2, with straight side walls, pointed corners, and a flat polar cap."
+ Velocities as fast as 400 km s are seen., Velocities as fast as –400 km $^{-1}$ are seen.
+ An interesting feature at the NEL position is the pair of bright spots that occupy. the equator of the LII. giving the impression of a slice through a Vilted equatorial ring (also seen in channel maps of Fe i1] AA4SO1 presented by Ishibashi et al..," An interesting feature at the NE1 position is the pair of bright spots that occupy the equator of the LH, giving the impression of a slice through a tilted equatorial ring (also seen in channel maps of [Fe ] $\lambda$ 4891 presented by Ishibashi et al.,"
+ and in gasdvnanical simulations of the LII: Gonzalez et 22004)., and in gasdynamical simulations of the LH; Gonzalez et 2004).
+ These. are the kinematic counterparts of the Purple Haze and 5 ΠΗ and N 11] features in images. which show marked temporal variability (Smith et 22000: 20042).," These are the kinematic counterparts of the Purple Haze and [S ] and [N ] features in images, which show marked temporal variability (Smith et 2000; 2004a)."
+ At the display. scale Ποσο here (see $4). a line drawn through these two features would trace out an equatorial plane tilted from the plane of the skv by ~40°. consistent with the inclination of the llomunculus (Smith 2002b: Davidson ct 22001).2e:," At the display scale chosen here (see 4), a line drawn through these two features would trace out an equatorial plane tilted from the plane of the sky by $\sim$ $\arcdeg$, consistent with the inclination of the Homunculus (Smith 2002b; Davidson et 2001).:"
+ Both polar lobes of the LU are clearly seen at this position. indicating that here we can see to the far side of the LII (Smith et 110998: Smith Gohrz 2000: Davidson et 22001: Smith 2002b).," Both polar lobes of the LH are clearly seen at this position, indicating that here we can see to the far side of the LH (Smith et 1998; Smith Gehrz 2000; Davidson et 2001; Smith 2002b)."
+ Interestingly. this is also the onlv slit. position exhibiting blueshifted equatorial Fe 11] emission from the “skirt” of the loqThe saturated stellar continuum. makes it. sinipossible. to measure the precise velocity of the LII along the line of sight. but considering the kinematic structure in adjacent slit positions. the blueshifted wall of the LII crosses our line ofsight tothestarat 140220 kms. + (shown with the white dot).," Interestingly, this is also the only slit position exhibiting blueshifted equatorial [Fe ] emission from the “skirt” of the The saturated stellar continuum makes it impossible to measure the precise velocity of the LH along the line of sight, but considering the kinematic structure in adjacent slit positions, the blueshifted wall of the LH crosses our line of sight to the star at $\pm$ 20 km $^{-1}$ (shown with the white dot)."
+ This agrees with the UV absorption component at 1H6 km + (Gull et 22004). confirming that the absorption feature is indeed from the LII.," This agrees with the UV absorption component at –146 km $^{-1}$ (Gull et 2004), confirming that the absorption feature is indeed from the LH."
+ The bright Fe 1] emission just northwest of the star at 46 km (Smith 2004) is from the Weigelt knots (Weigelt LEhersberecr 1986).2d):, The bright [Fe ] emission just northwest of the star at –46 km $^{-1}$ (Smith 2004) is from the Weigelt knots (Weigelt Ebersberger 1986).:
+ The SWI position gives the best representation of the bipolar structure of the LII. with two nearly complete. elosed polar lobes.," The SW1 position gives the best representation of the bipolar structure of the LH, with two nearly complete, closed polar lobes."
+ It exhibits the samo ILat-topped or trapezoidal structure as is seen at NET., It exhibits the same flat-topped or trapezoidal structure as is seen at NE1.
+ However. SWI begins to show asymmetry in the LII: the redshifted NW lobe is about Larger than the SE lobe.," However, SW1 begins to show asymmetry in the LH; the redshifted NW lobe is about larger than the SE lobe."
+ Also. unlike NEL. there is a strong brightness asymmetry in the emission knots associated: with the equatorial ring (the blueshifted knot is much brighter. as for the slit passing through the star).2," Also, unlike NE1, there is a strong brightness asymmetry in the emission knots associated with the equatorial ring (the blueshifted knot is much brighter, as for the slit passing through the star).:"
+ek Structure at. this. position. further accentuates the asvmmetry of the LIL., Structure at this position further accentuates the asymmetry of the LH.
+ Phe blueshitted lobe is much smaller than its counterpart at NEP. and it is about half the size of its own redshifted lobe.," The blueshifted lobe is much smaller than its counterpart at NE2, and it is about half the size of its own redshifted lobe."
+ Obscuration of the redshifted lobe by dust in the equator is apparent., Obscuration of the redshifted lobe by dust in the equator is apparent.
+ The shape of the LII and the way its kinematic structure is clisplaved in Figure 2 can give valuable clues. to its age., The shape of the LH and the way its kinematic structure is displayed in Figure 2 can give valuable clues to its age.
+ The relative scale between the spatial ancl velocity directions in Figure 2 reflects the age. in the sense that a horizontal stretch implies vounger material. ancl horizontal compression implies older material.," The relative scale between the spatial and velocity directions in Figure 2 reflects the age, in the sense that a horizontal stretch implies younger material, and horizontal compression implies older material."
+ Since the LII has some inherent asvmmetry. the choice of the horizontal stretch is subjective. depending on which features one uses to gauge the appropriate scaling.," Since the LH has some inherent asymmetry, the choice of the horizontal stretch is subjective, depending on which features one uses to gauge the appropriate scaling."
+ Figure 2 is clisplaved with a scale of πμ πο 1. Corresponding. to an ejection⊲⋠ date around 1910.," Figure 2 is displayed with a scale of $\arcsec$ =117 km $^{-1}$, corresponding to an ejection date around 1910."
+ This was chosen to match proper motions of the Weigelt knots (Smith et 22004a)., This was chosen to match proper motions of the Weigelt knots (Smith et 2004a).
+ Also. at this scale. a line drawn through the two bright equatorial knots at the NET position is tilted from vertical by 40. which matches the inclination of the Homunculus. as noted above.," Also, at this scale, a line drawn through the two bright equatorial knots at the NE1 position is tilted from vertical by $\sim$ $\arcdeg$, which matches the inclination of the Homunculus, as noted above."
+ llowever. at this cdisplav scale. some portions of the LIE still appear somewhat stretched. which has two ramifications.," However, at this display scale, some portions of the LH still appear somewhat stretched, which has two ramifications."
+ First. if one assumes axial svmmetry. it confirms that the Llomunculus was not ejected during the Great Eruption in the 15105. because such features would »' horizontally compressed. rather than elongated.," First, if one assumes axial symmetry, it confirms that the Homunculus was not ejected during the Great Eruption in the 1840s, because such features would be horizontally compressed, rather than elongated."
+ Second. he polar features look the most svmametric at a display scale corresponding to a later ejection date of 1920.1930.," Second, the polar features look the most symmetric at a display scale corresponding to a later ejection date of 1920–1930."
+ The potential reconciliation of the age discrepancy may jwe to do with acceleration of ejecta. as suggested already w Smith et ((2004b).," The potential reconciliation of the age discrepancy may have to do with acceleration of ejecta, as suggested already by Smith et (2004b)."
+ Material ejected in the 1890 event hat has been accelerated by radiation pressure or stellar wind ram pressure would show Laster Doppler shifts and ugher proper motion than expected., Material ejected in the 1890 event that has been accelerated by radiation pressure or stellar wind ram pressure would show faster Doppler shifts and higher proper motion than expected.
+ As noted: already. equatorial zones of the LIL have probably. been accelerated since an 1890 ejection (Smith et 22004b).," As noted already, equatorial zones of the LH have probably been accelerated since an 1890 ejection (Smith et 2004b)."
+ In. Figure 2 he polar features appear elongated: even for an ejection date of 1910 if they originated. in the 1800. event as well. they must have been accelerated. even. more than the corresponding equatorial features.," In Figure 2 the polar features appear elongated even for an ejection date of 1910 — if they originated in the 1890 event as well, they must have been accelerated even more than the corresponding equatorial features."
+ La polar cirections. ram pressure of the wind probably dominates. since η Car has à latitucle-cdependent wind with an cllective mass-loss rate of Al10 CAM. ! and polar speeds of 600 km s.1 (Smith et 22003a).," In polar directions, ram pressure of the wind probably dominates, since $\eta$ Car has a latitude-dependent wind with an effective mass-loss rate of $\dot{M}\simeq$ $^{-3} M_{\odot}$ $^{-1}$ and polar speeds of 600 km $^{-1}$ (Smith et 2003a)."
+ Thus. the polar wind speed is faster than the LII. so we should. expect some interaction.," Thus, the polar wind speed is faster than the LH, so we should expect some interaction."
+ The [act that the polar caps of the LII have so far maintained their integrity means that the fast stellar wind has not been able to plow through the LIE or disrupt it through, The fact that the polar caps of the LH have so far maintained their integrity means that the fast stellar wind has not been able to plow through the LH or disrupt it through
+curve. and was devised by. Morbidelli ct al (1997).,"curve, and was devised by Morbidelli et al (1997)."
+ From the observed. Hi [Duxes. the colour excesses E(AV7) have been found using the intrinsic colours by Wegner (1994).," From the observed IR fluxes, the colour excesses $(\lambda - V)$ have been found using the intrinsic colours by Wegner (1994)."
+ The quantity ly has then been determined with a least squares solution. by fitting (Cardelli et al. 1959): where A=R.1.J.Kk.L and B is the extinction curve ux taken from Rieke Lebolsky From the derived cl) and the known ECGDΕ owe obtain #y.," The quantity $A_V$ has then been determined with a least squares solution, by fitting (Cardelli et al, 1989): where $\lambda = R,I,J,K,L$ and $R_L$ is the extinction curve $\frac{A_\lambda}{A_V}$ taken from Rieke Lebofsky From the derived $A_V$ and the known $(B-V)$ we obtain $R_V$."
+ Since the fittinge equation 3 is a homogeneousὃν one. the uncertainty of each shy hase been computed. by considering that we have NVl degree of freedom. (GN being the number of photometric bands available).," Since the fitting equation \ref{morbi} is a homogeneous one, the uncertainty of each $A_V$ hase been computed by considering that we have $N-1$ degree of freedom $N$ being the number of photometric bands available)."
+ Εις implies that we can obtain only a lower limit on the 74 uncertainty. since it is not easy to take into account spectral mis-classilication anc hence inaccuracy in the adopted intrinsic colors (Patriarchi οἱ al The results are summarized. in Table 3. where the identification is reported together with the individual reddening E(5Εν the total absorption zl. and the ratio of total-to-selective absorption £&_=Enwet: these«o latter derived from the method.," This implies that we can obtain only a lower limit on the $R_V$ uncertainty, since it is not easy to take into account spectral mis-classification and hence inaccuracy in the adopted intrinsic colors (Patriarchi et al The results are summarized in Table 3, where the identification is reported together with the individual reddening $(B-V)$, the total absorption $A_V$, and the ratio of total-to-selective absorption $R_V~=~\frac{A_V}{E(B-V)}$, these latter derived from the method."
+(0). Phe last column reports the value of Ay derived from the methodfa., The last column reports the value of $R_V$ derived from the method.
+ A weighted mean vields Ay=3.09+044., A weighted mean yields $R_V~=~3.09\pm0.44$.
+ in fine agreement both with the results of method fa). and with the finding of the dilferential reddening method.," in fine agreement both with the results of method , and with the finding of the differential reddening method."
+ In [ασ the three values we obtain [or Ry clearly overlap. and a weighted: mean vields the final estimate which means that the interstellar reddening law toward ‘Trumpler 15 appears normal. as previously suggested: for instance bv Feinsten et al (1980).," In fact the three values we obtain for $R_V$ clearly overlap, and a weighted mean yields the final estimate which means that the interstellar reddening law toward Trumpler 15 appears normal, as previously suggested for instance by Feinsten et al (1980)."
+ We already obtained. an estimate of Trumpler 15 distance ron the analysis of the variable extinction plot in Section 6., We already obtained an estimate of Trumpler 15 distance from the analysis of the variable extinction plot in Section 6.
+ Lere we address the issues of the distance. ancl the age clirectly by considering the reddening corrected CAIDs for he member stars., Here we address the issues of the distance and the age directly by considering the reddening corrected CMDs for the member stars.
+" The correction to the magnitudes is done w computing V,—Vo300V) for each star.", The correction to the magnitudes is done by computing $V_o = V -3.0 \times E(B-V)$ for each star.
+ This is plausible since we have shown that the reddening law in he direction of the clusteris normal., This is plausible since we have shown that the reddening law in the direction of the clusteris normal.
+complete circle ancl being thus called the Sextant (Efremov Elmeereen 1998).,complete circle and being thus called the Sextant (Efremov Elmegreen 1998).
+ Efremov. Ehlerova Palous (1999) garowed that the OB associations are regularly spaced. 10 deprojected average distance between two subsequent stellar groups being ~ 37 pe.," Efremov, Ehlerova Palous (1999) showed that the OB associations are regularly spaced, the deprojected average distance between two subsequent stellar groups being $\sim$ 37 pc."
+ Furthermore. using numerical simulations of shells expanding in a mass model of the Large Magellanie Cloud. they concluded that the formation of these 5 OB associations is most probably the result of a triggered star formation episode in the supershell created by ολο located near the Sextant centre.," Furthermore, using numerical simulations of shells expanding in a mass model of the Large Magellanic Cloud, they concluded that the formation of these 5 OB associations is most probably the result of a triggered star formation episode in the supershell created by SNeII located near the Sextant centre."
+ Pheir simulations also explain how projection ellects make the visible star forming regions only a fraction of a total circle., Their simulations also explain how projection effects make the visible star forming regions only a fraction of a total circle.
+ The influence of the number of clumps embedded within the shell on the transverse collapse is. illustrated in Fig. 6.., The influence of the number of clumps embedded within the shell on the transverse collapse is illustrated in Fig. \ref{fig:disc_eta}.
+ The parameter is not as negligeable as the initial perturbed surface density., The parameter $\eta$ is not as negligeable as the initial perturbed surface density.
+ In some cases. it acts upon the fragmentation process almost as strongly as the initial transverse velocity.," In some cases, it acts upon the fragmentation process almost as strongly as the initial transverse velocity."
+ Lt is thus worth estimating the angular wavenumbers which are the most favourable to the growth of an initial perturbation., It is thus worth estimating the angular wavenumbers which are the most favourable to the growth of an initial perturbation.
+ In order to do so. we now derive an analvtical approximation of the instantaneous growth rate of the perturbation.," In order to do so, we now derive an analytical approximation of the instantaneous growth rate of the perturbation."
+ This is done straightforwardly by assuming that the perturbecl quantities vary exponentially. with time /. in analogy with the exponential growth. rate found. in other instability problems (e.g. the Jeans. mass) even though there is no real exponential growth in this problem owing to the shell expansion and the corresponding time dependence of ay.," This is done straightforwardly by assuming that the perturbed quantities vary exponentially with time $t$, in analogy with the exponential growth rate found in other instability problems (e.g. the Jeans mass) even though there is no real exponential growth in this problem owing to the shell expansion and the corresponding time dependence of $\sigma _0$."
+ I is thus important to keep in mind that the perturbation growth rate derived. below (eqs., It is thus important to keep in mind that the perturbation growth rate derived below (Eqs.
+ 28 and 31)) is used to compute the temporal evolution of σι (the cases displaved in Fig.," \ref{eq:omega_eta}
+ and \ref{eq:fg_omega}) ) is used to compute the temporal evolution of $\sigma _1$ (the cases displayed in Fig."
+ 6 are obtained from solving Ίσα» 17--18)) but merely to derive 7 estimates favouring the shell collapse., \ref{fig:disc_eta} are obtained from solving Eqs \ref{eq:per_cont_num_phi}- \ref{eq:per_motion_num_phi}) ) but merely to derive $\eta$ estimates favouring the shell collapse.
+ Phe perturbed quantities are written: where w is the angular frequcney of the perturbation., The perturbed quantities are written: and where $\omega$ is the angular frequency of the perturbation.
+ Following Eqs., Following Eqs.
+ 23 and 24. we write w for the time derivatives and ηης for. the transverse. eradients.," \ref{sigma1_exp} and \ref{v_exp}, , we write $\omega$ for the time derivatives and $- i \eta/R_s$ for the transverse gradients."
+ Therefore. Eqs.," Therefore, Eqs."
+ 1: and 2. become respectively and using Eq.," \ref{eq:per_cont} and \ref{eq:per_motion}
+ become respectively and using Eq."
+ 3. in the latter., \ref{g1_sigma1} in the latter.
+ The elimination of the perturbecl quantities σι and ve between Eqs., The elimination of the perturbed quantities $\sigma _1$ and $v$ between Eqs.
+ 25. and 26 provides the dispersion equation. namely the relation between the angular frequency c (i.e. the instantaneous growth rate) and the angular wavenumber η o the perturbation: whose solution is given by Let us consicer the inodo.," \ref{eq:per_cont2} and \ref{eq:per_motion2} provides the dispersion equation, namely the relation between the angular frequency $\omega$ (i.e. the instantaneous growth rate) and the angular wavenumber $\eta$ of the perturbation: whose solution is given by Let us consider the ."
+ This one corresponds to the sequence of values of 5 which maximises the angular. frequency a at cach moment of the shell propagation. Le. ary=1) such that Equation 29— indeed. corresponds. to a maximum. since the discriminant of Eq.," This one corresponds to the sequence of values of $\eta$ which maximises the angular frequency $\omega$ at each moment of the shell propagation, i.e. $\eta _{fg} = \eta (t)$ such that Equation \ref{eq:fg_def} indeed corresponds to a maximum since the discriminant of Eq."
+ 28. shows a negative curvature with 4g., \ref{eq:omega_eta} shows a negative curvature with $\eta$ .
+" The instantaneous angular wavenumber and the instantaneous angular [frequency associated: to. the first growing mode obev respectively and ‘Taking into account the dependence of AZ and 7, (Eq. 9))", The instantaneous angular wavenumber and the instantaneous angular frequency associated to the first growing mode obey respectively and Taking into account the dependence of $M$ and $R_s$ (Eq. \ref{eq:RsHPBt1/3}) )
+" on £, and IN. Eq."," on $P_h$ and $N$, Eq."
+ 30. shows that yyy depends on the external pressure P. on the SN number NV and on time fas A first guess of a favourable angular wavenumber can be estimated from the temporal average of Eq.," \ref{eq:fg_eta} shows that $\eta _{fg}$ depends on the external pressure $P_h$, on the SN number $N$ and on time $t$ as A first guess of a favourable angular wavenumber can be estimated from the temporal average of Eq."
+ 30. over the time spent by the supershell in the hot protogalactic background: where Af ds. the duration. of the SN. phase., \ref{eq:fg_eta} over the time spent by the supershell in the hot protogalactic background: where $\Delta t$ is the duration of the SN phase.
+ Among the cases cdisplaved in Fig., Among the cases displayed in Fig.
+" 6 (e,=1 Alyy)=0.01690,). OO)=0.01.0, ). we see that the shell transverse collapse is achieved. if. for instance. I4=10 Ü'addeneeem7. N=100 and 4216 (top panel) or if [4=510J'ddyenecnm 7. N=LOO and g-10 (middle panel)."," \ref{fig:disc_eta} $c_s$ $^{-1}$, $\tilde \sigma _1(t_{em})=0.01\sigma _0(t_{em})$, $v(t_{em})=0.01\,V_s(t_{em})$ ), we see that the shell transverse collapse is achieved if, for instance, $P_h = 10^{-10}$ $^{-2}$, $N$ =100 and $\eta$ =16 (top panel) or if $P_h = 5 \times 10^{-10}$ $^{-2}$, $N$ =100 and $\eta$ =10 (middle panel)."
+ It is interesting to note that these values of a reasonably match those given. by Ίσα. 33... ," It is interesting to note that these values of $\eta$ reasonably match those given by Eq. \ref{eta_fg_ave}, ,"
+"for the above mentioned cases. Le.. «€gygy 7-12 and <gy, 7-9. respectively,"," for the above mentioned cases, i.e., $<\eta _{fg}>$ =12 and $<\eta _{fg}>$ =9, respectively."
+ As Eq., As Eq.
+ 33. has been derived. under the assumption of an exponential growth rate. this agreement a posteriori justifies itsvalidity as a convenient estimate of the number of forming clumps in collapsing shells.," \ref{eta_fg_ave}
+ has been derived under the assumption of an exponential growth rate, this agreement a posteriori justifies itsvalidity as a convenient estimate of the number of forming clumps in collapsing shells."
+ The development of a gravitational instability within a supershell also depends on the sound. speed e; of the shell material., The development of a gravitational instability within a supershell also depends on the sound speed $c_s$ of the shell material.
+" Indeed. ος is cirecthy related to the thermal pressure £, of the shell gas through In this equation. & and mg are the Boltzmann constant andthe hwdrogen mass. respectively. while ps. Zi and fy are the mass density. the temperature andthe mean molecular weightof the shell.respectively."," Indeed, $c_s$ is directly related to the thermal pressure $P_s$ of the shell gas through In this equation, $k$ and $m_H$ are the Boltzmann constant andthe hydrogen mass, respectively, while $\rho _s$ , $T_s$ and $\mu _s$ are the mass density, the temperature andthe mean molecular weightof the shell,respectively."
+ The larger the, The larger the
+o surface brightnesses can more clearly show deviations in he relation between two wave bands. including systematic variations in colour related to surface brightness.,"to surface brightnesses can more clearly show deviations in the relation between two wave bands, including systematic variations in colour related to surface brightness."
+ Lf à onc-o-one correspondence exists between the PALL and 24 jun surface brightnesses. then the slopes of the best fit lines in hese plots would. be close to 0. and the scatter in. these ots would be small.," If a one-to-one correspondence exists between the PAH and 24 $\mu$ m surface brightnesses, then the slopes of the best fit lines in these plots would be close to 0, and the scatter in these plots would be small."
+ Such plots would be biased: towards »oducineg relations with slopes of -1 in log-log space if the wo wave bands were randomly distributed., Such plots would be biased towards producing relations with slopes of -1 in log-log space if the two wave bands were randomly distributed.
+ Llowever. since he Εαν densities in all wave bands. we studied: here are approximately proportional to each other. such biases will not be present.," However, since the flux densities in all wave bands we studied here are approximately proportional to each other, such biases will not be present."
+ To measure the physical (intrinsic) scatter around the best fit line. we will subtract the sum of the square of the measurement uncertainties from the sum. of the square of the residuals from the best fit line.," To measure the physical (intrinsic) scatter around the best fit line, we will subtract the sum of the square of the measurement uncertainties from the sum of the square of the residuals from the best fit line."
+" “Phis is eiven by where ας and y are abscissa ancl ordinate valucs with corresponding uncertainties m, and m, and e ancl b are the v-intereept and slope of of the best fit line.", This is given by where $x$ and $y$ are abscissa and ordinate values with corresponding uncertainties $\sigma_x$ and $\sigma_y$ and $a$ and $b$ are the y-intercept and slope of of the best fit line.
+ A value of 0 is reported if the result from Equation 10. is negative. as this would indicate that all of the scatter around the best fit line could be accounted for by the measurement. uncertainties.," A value of 0 is reported if the result from Equation \ref{e_scatter} is negative, as this would indicate that all of the scatter around the best fit line could be accounted for by the measurement uncertainties."
+ Figure 2. shows how the (PALL δ j/m)/24 jun ratio varies with 24 jam surface brightness among the 45 arcsec square regions described in Section. 2.4.., Figure \ref{f_pahvs24} shows how the (PAH 8 $\mu$ m)/24 $\mu$ m ratio varies with 24 $\mu$ m surface brightness among the 45 arcsec square regions described in Section \ref{s_dataprep}.
+ Note that. the resolution of the data used in this figure is matched. to the 38 aresee resolution of the 160 jm images so that the results from these figures can be more easily compared. to the results in Section 4.., Note that the resolution of the data used in this figure is matched to the 38 arcsec resolution of the 160 $\mu$ m images so that the results from these figures can be more easily compared to the results in Section \ref{s_comp_pah160}.
+ Phe slopes and intrinsic scatter for the best fit lines are given in Table 3.., The slopes and intrinsic scatter for the best fit lines are given in Table \ref{t_pahvs24}.
+ The best fitting lines are determined. using uncertainties in both the x- and vedirections to weight the data. so the fits are strongly weighted towards high surface brightness regions.," The best fitting lines are determined using uncertainties in both the x- and y-directions to weight the data, so the fits are strongly weighted towards high surface brightness regions."
+ The slopes ofthe best fit lines in Figure 2. are generally not statistically equivalent to 0. which demonstrates. that the (PALL S p/m)/24 jm ratio varies with surface brightness.," The slopes of the best fit lines in Figure \ref{f_pahvs24} are generally not statistically equivalent to 0, which demonstrates that the (PAH 8 $\mu$ m)/24 $\mu$ m ratio varies with surface brightness."
+ ln many galaxies (most notably NGC 3031. NGC 3351. and NGC 6946). the (PALL S jm)/24 jn ratio decreases notably as the 24 yam surface brightness increases.," In many galaxies (most notably NGC 3031, NGC 3351, and NGC 6946), the (PAH 8 $\mu$ m)/24 $\mu$ m ratio decreases notably as the 24 $\mu$ m surface brightness increases."
+ These tend to be galaxies with infrared-bright point-like nuclei., These tend to be galaxies with infrared-bright point-like nuclei.
+ In a lew other galaxies (NGC 2403. and NGC 7793. [or example). the (PALL δ sam)/24 jm ratio increases as the 24 pm surface brightness increases.," In a few other galaxies (NGC 2403 and NGC 7793, for example), the (PAH 8 $\mu$ m)/24 $\mu$ m ratio increases as the 24 $\mu$ m surface brightness increases."
+ The data for some galaxics also show a broad scatter. particularly for galaxics without infrared-bright nuclei.," The data for some galaxies also show a broad scatter, particularly for galaxies without infrared-bright nuclei."
+ For a given 24 yam surface brightness. the (PALL S sam)/24 pim ratio varies by over a factor of 2 in high signal-to-noise regions in NGC 925 and NGC 2408.," For a given 24 $\mu$ m surface brightness, the (PAH 8 $\mu$ m)/24 $\mu$ m ratio varies by over a factor of 2 in high signal-to-noise regions in NGC 925 and NGC 2403."
+ Although both of these galaxies are fit with lincar relations in Figure 2.. the broad. scatter indicates that such a fit is unrealistic. so the (PALL 8) and 24 jim emission must be only weakly associated with cach other.," Although both of these galaxies are fit with linear relations in Figure \ref{f_pahvs24}, the broad scatter indicates that such a fit is unrealistic, so the (PAH 8) and 24 $\mu$ m emission must be only weakly associated with each other."
+ Also. some data points fall well below the best fit. lines in Figure 2.. xuwiicularlv in the plots for NGC 2403. and NGC 3938.," Also, some data points fall well below the best fit lines in Figure \ref{f_pahvs24}, particularly in the plots for NGC 2403 and NGC 3938."
+ These regions correspond to very infrared-bright star-forming regions visible in Figure 1.., These regions correspond to very infrared-bright star-forming regions visible in Figure \ref{f_map}.
+ ligure 1. does not reveal the presence of any obvious dependence of the (PALE S jm)/24 pm on radius. but it is still useful to measure such gradients for comparison with abundance gradients. especially since it has been shown that the (PALL S μι)δε pim ratio varies with metallicity," Figure \ref{f_map} does not reveal the presence of any obvious dependence of the (PAH 8 $\mu$ m)/24 $\mu$ m on radius, but it is still useful to measure such gradients for comparison with abundance gradients, especially since it has been shown that the (PAH 8 $\mu$ m)/24 $\mu$ m ratio varies with metallicity"
+where Combining this equation with eq. (44)),where Combining this equation with eq. \ref{im}) )
+ gives Solving for & gives the critical wavenumber & above which the svstem is unstable: independent. of entrainment., gives Solving for $k$ gives the critical wavenumber $k_c$ above which the system is unstable: independent of entrainment.
+ The critical wavenumber ἂν. is minimized for 06=tan1/2)., The critical wavenumber $k_c$ is minimized for $\theta=\tan^{-1}(\sqrt{2})$.
+" Por Kcbe. we have the approximate solutions For the outer core. e,23«10.7 forc, =0.05 (?).."," For $k>>k_c$, we have the approximate solutions For the outer core, $\epsilon_n\simeq 3\times 10^{-3}$ for $x_p=0.05$ \citep{ch06}."
+ Though entrainment is essential in producing vortex. pinning. it has a negligibleBUS ellect on the 5growth rate of the instability. and. no ellect on the critical wavenumber. so we ignore5 it in the estimates below.," Though entrainment is essential in producing vortex pinning, it has a negligible effect on the growth rate of the instability, and no effect on the critical wavenumber, so we ignore it in the estimates below."
+" The instability arises from coupling between the velocity difference w,,,, and the neutron vorticity ων through the force of eq. 26].", The instability arises from coupling between the velocity difference $\wbf_{np}$ and the neutron vorticity $\omegabf_n$ through the force of eq. \ref{df}] ].
+ Dissipation damps perturbations for &«fy. but for &cA the finite vortex mobility gives rise to growing perturbations. under the Magnus force.," Dissipation damps perturbations for $k<k_c$, but for $k>k_c$ the finite vortex mobility gives rise to growing perturbations under the Magnus force."
+H For &7Ay. the growth rate scales as (àAeg)2F7.," For $k>>k_c$, the growth rate scales as $(\alpha\beta w_0)^{1/2}$."
+ For d««os df takes a constant value. but the erowth rate of the mode becomes small. going to zero as 3 goes to zero.," For $\beta<<\alpha$, $k_c$ takes a constant value, but the growth rate of the mode becomes small, going to zero as $\beta$ goes to zero."
+ In the highlv-damped regime. 37a. damping restricts the unstable mode to large &. generally stabilizing the system.," In the highly-damped regime, $\beta>>\alpha$ , damping restricts the unstable mode to large $k$, generally stabilizing the system."
+" There are no unstable modes for either a=0 or 3—0: the instability occurs only if vortex motion has components along both di, and (icWap)owy.", There are no unstable modes for either $\alpha=0$ or $\beta=0$; the instability occurs only if vortex motion has components along both $\hat{w}_{np}$ and $(\hat{\omega}_n\times\hat{w}_{np})\times\hat{\omega}_n$.
+ We now confirm that the viscous stress is negligible for the @ unstable mode., We now confirm that the viscous stress is negligible for the $\sigma_-$ unstable mode.
+ The magnetic stress force per unit volume on the charged [uid [rom eq.(25)) is while the viscous force is Fourier transforming.e and usingὃν the induction equationi (28)). ogives ‘This ratio is largest for high wavenumber.," The magnetic stress force per unit volume on the charged fluid from \ref{paccel}) ) is while the viscous force is Fourier transforming, and using the induction equation \ref{induction}) ), gives This ratio is largest for high wavenumber."
+" In this limit In the outer core. £4 is typically 10"" emo ΕΦ}."," In this limit In the outer core, $\nu_e$ is typically $\sim 10^6$ $^2$ $^{-1}$ \citep{cl87,acg05}."
+ Below we estimate ad~10.D., Below we estimate $\alpha\beta\sim 10^{-18}$.
+" The hydrodynamic treatment requires hep<<Qu. so 8f,fofi,107 for a typical neutron star with Qu=100. aand By=1077 Gt the viscous force is negligible compared to the magnetic force Lor these Iow-[requeney. modes."," The hydrodynamic treatment requires $kc_T<<\Omega_0$, so $\delta f_v/\delta f_m\lap
+10^{-9}$ for a typical neutron star with $\Omega_0=100$ and $B_0=10^{12}$ G; the viscous force is negligible compared to the magnetic force for these low-frequency modes."
+ ‘To obtain the growth rate of the instability. we now estimate the pinning parameters o and 2 for the vortex creep process.," To obtain the growth rate of the instability, we now estimate the pinning parameters $\alpha$ and $\beta$ for the vortex creep process."
+ To make these estimates. we regard the process of vortex creep as consisting of two clistinet states of motion for a given vortex segment.," To make these estimates, we regard the process of vortex creep as consisting of two distinct states of motion for a given vortex segment."
+ Most. of the time. the vortex segment is pinned.," Most of the time, the vortex segment is pinned."
+ A small fraction of the time. the vortex segment is translating against a drag force to a new pinning configuration.," A small fraction of the time, the vortex segment is translating against a drag force to a new pinning configuration."
+ The mutual friction force we are using (eq. 21)), The mutual friction force we are using (eq. \ref{fmf}) )
+ is This force represents the spatially and temporally averaged force exerted on the neutron Iuid by the creep process., is This force represents the spatially and temporally averaged force exerted on the neutron fluid by the creep process.
+ The first term is the Magnus force for perfect pinning. while the remaining terms give the contribution to the force due to vortex motion.," The first term is the Magnus force for perfect pinning, while the remaining terms give the contribution to the force due to vortex motion."
+ Whena vortex segment is unpinned and moving against drag. we take the force to have the same form. but with different cocllicients:," Whena vortex segment is unpinned and moving against drag, we take the force to have the same form, but with different coefficients:"
+survey of the DIBs in a large sample of stars.,survey of the DIBs in a large sample of stars.
+ High resolution (R~ 37.500) visible (4000-10000 Aj) spectra have been obtained with the Astrophysical Research Consortium Echelle Spectrograph (ARCES) on the 3.5-m telescope at the Apache Point Observatory.," High resolution $R \sim$ 37,500) visible (4000–10000 ) spectra have been obtained with the Astrophysical Research Consortium Echelle Spectrograph (ARCES) on the 3.5-m telescope at the Apache Point Observatory."
+ Data reduction is performed using standard IRAF routines. as described in detail by Thorburn(2000).," Data reduction is performed using standard IRAF routines, as described in detail by \citet{thorburn}."
+. A more complete deseription of our DIB survey will be given in à future paper., A more complete description of our DIB survey will be given in a future paper.
+" Given the constants from Lakinetal.(2000) (872897 MHz. B'z887 MHz. AQ,227.4!.. and At,20.6 em). we used the method of HillandVanVleck(1928) to calculate the energy levels of aand the intensity. factors for the individual rotational. lines within a given vibronic band. ["," Given the constants from \citet{lakin}
+ $B''$ =897 MHz, $B'$ =887 MHz, $A_{SO}''$ =27.4, and $A_{SO}'$ =0.6 ), we used the method of \citet{vanvleck} to calculate the energy levels of and the intensity factors for the individual rotational lines within a given vibronic band. ["
+We assumed the same constants for each vibronic band. as the vibrational dependence of the constants Is expected to be smaller than the uncertainty in the determined constants.],"We assumed the same constants for each vibronic band, as the vibrational dependence of the constants is expected to be smaller than the uncertainty in the determined constants.]"
+ The populations of the rotational states of wwere then calculated using a Boltzmann expression at a given temperature. from which we were able to simulate the absorption spectrum.," The populations of the rotational states of were then calculated using a Boltzmann expression at a given temperature, from which we were able to simulate the absorption spectrum."
+ Temperatures between 10 K and 90 K were considered. as diffuse clouds can be expected to have temperatures within this range.," Temperatures between 10 K and 90 K were considered, as diffuse clouds can be expected to have temperatures within this range."
+ For the linewidth of each transition. we assumed a Gaussian profile with full-width at half-maximum 10 km/s. which is the FWHM of the observed K lines in HD 185418 and HD 229059 (two stars we have chosen for the comparsion due to their narrow K lines).," For the linewidth of each transition, we assumed a Gaussian profile with full-width at half-maximum 10 km/s, which is the FWHM of the observed K lines in HD 185418 and HD 229059 (two stars we have chosen for the comparsion due to their narrow K lines)."
+ We begin by considering the Q”=1/2 spin-orbit component of the origin (00) band ofCz.. in comparison with the 46270 DIB.," We begin by considering the $\Omega''$ =1/2 spin-orbit component of the origin $0^0_0$ ) band of, in comparison with the $\lambda$ 6270 DIB."
+ The origin band is naturally the strongest of the laboratory features. and A6270 is also by far the strongest of the DIBs suggested to correspond toC5.," The origin band is naturally the strongest of the laboratory features, and $\lambda$ 6270 is also by far the strongest of the DIBs suggested to correspond to."
+. Figure | shows the spectra of A6270 toward HD 185418 and HD 229059. along with the simulations of the oorigin band at temperatures of 10. 30. 50. and 70 K. As can be seen from the figure. neither the central wavelengths nor the profiles of the spectra agree with the 46270 diffuse band.," Figure \ref{origin} shows the spectra of $\lambda$ 6270 toward HD 185418 and HD 229059, along with the simulations of the origin band at temperatures of 10, 30, 50, and 70 K. As can be seen from the figure, neither the central wavelengths nor the profiles of the spectra agree with the $\lambda$ 6270 diffuse band."
+ This disagreement argues strongly against the assignment of A6270 toC5., This disagreement argues strongly against the assignment of $\lambda$ 6270 to.
+". In Figure 2. we consider both the O""—-1/2 (left) and 0723/2 (right) components of the oorigin band.", In Figure \ref{originsat} we consider both the $\Omega''$ =1/2 (left) and $\Omega''$ =3/2 (right) components of the origin band.
+ Because Q/=3/2 is higher in energy. the intensity of the right-hand component increases with temperature. as evident in the simulations at 30. 60. and 90 K. In Figure 2.. an (unreasonably large) Gaussian linewidth of 30 km/s has been assumed in order to improve the agreement with A6270.," Because $\Omega''$ =3/2 is higher in energy, the intensity of the right-hand component increases with temperature, as evident in the simulations at 30, 60, and 90 K. In Figure \ref{originsat}, an (unreasonably large) Gaussian linewidth of 30 km/s has been assumed in order to improve the agreement with $\lambda$ 6270."
+" It is difficult to state with certainty because of the presence of the strong A6284 DIB. but it appears that there is little evidence for the 0""23/2 component in the astronomical spectra."," It is difficult to state with certainty because of the presence of the strong $\lambda$ 6284 DIB, but it appears that there is little evidence for the $\Omega''$ =3/2 component in the astronomical spectra."
+" However. because the intrinsic profile of the strong A6284 DIB ts not known. the presence of the O""—-3/2 component cannot be definitively ruled out."," However, because the intrinsic profile of the strong $\lambda$ 6284 DIB is not known, the presence of the $\Omega''$ =3/2 component cannot be definitively ruled out."
+galaxies.,galaxies.
+ Lhe most important selection criteria are (1) ασ(2): the fraction of total flux contributed by the ce Vaucouleurs component when the galaxy. profile is fitted. by the sum of de Vaucouleurs ancl exponential profiles., The most important selection criteria are (i) $ \rm frac_{deV}(r)$: the fraction of total flux contributed by the de Vaucouleurs component when the galaxy profile is fitted by the sum of de Vaucouleurs and exponential profiles.
+ We include only non-cdisk galaxies with fracqa of unity lor both the g and the x band. (, We include only non-disk galaxies with $\rm frac_{deV}$ of unity for both the $g$ and the $r$ band. (
+ii)Eclass. a 1D classification of spectral tvpo. obtained from a principal component analvsis (Yip et al.,"ii), a 1D classification of spectral type, obtained from a principal component analysis (Yip et al."
+ 2004) lower (negative) values corresponding to absorption-line galaxies with old. stellar populations. positive values indicating star-forming galaxies.," 2004) – lower (negative) values corresponding to absorption-line galaxies with old stellar populations, positive values indicating star-forming galaxies."
+" We include only galaxies with eclass«0. In addition we require (ii) Concentration. defined as the ratio of the Petrosian radius ρου (he angular radius containing 90% of the Petrosian flux) to the smaller radius rp;,50. (containing 50%). measured. in the ;band. must. be 2.1 to reduce possible contamination of later-tvpe galaxies (this selection excludes onlv a few (iv) The r band magnitude. corrected. for Galactic reddening. must be 14.5<r«17.5 (this excludes some of the faintest (v) The velocity dispersion (0) as measured. from the SDSS spectrum must be at least 60 kni s|. but. no more than 450 km s1 (the upper limitD. to exclude superimposed. galaxies. as no single galaxy in the SDSS appears to have -450 km s.|: eg Dernardi et al."," We include only galaxies with $\rm {\tt eclass}<0$ In addition we require (iii) Concentration, defined as the ratio of the Petrosian radius $r_{Pet, 90}$ (the angular radius containing $90\%$ of the Petrosian flux) to the smaller radius $r_{Pet,50}$ (containing $50\%$ ), measured in the $i$ -band, must be $>2.4$ to reduce possible contamination of later-type galaxies (this selection excludes only a few (iv) The $r$ band magnitude, corrected for Galactic reddening, must be $14.5<r<17.5$ (this excludes some of the faintest (v) The velocity dispersion $\sigma$ ) as measured from the SDSS spectrum must be at least 60 km $\rm s^{-1}$, but no more than 450 km $\rm s^{-1}$ (the upper limit to exclude superimposed galaxies, as no single galaxy in the SDSS appears to have $\sigma>450$ km $s^{-1}$; e.g Bernardi et al."
+ 2008: Salviander et al. (, 2008; Salviander et al. (
+vi) Redshift must be 0.005—z<0.36. The reason for the upper limit is that. at 2> 0.36. the rest-frame r-band begins to shift significantly outside of the 0200. range of the SDSS spectra. so we cannot obtain an accurate rbanel k-correction.,"vi) Redshift must be $0.005\leq z \leq 0.36$ The reason for the upper limit is that, at $z>0.36$ , the rest-frame $r$ -band begins to shift significantly outside of the $\rm 9200\AA$ range of the SDSS spectra, so we cannot obtain an accurate $r$ -band k-correction."
+ The selected galaxies. satisfving all the above. number TOSTS.," The selected galaxies, satisfying all the above, number 70378."
+ We apply a number of. generally very small. corrections to the σ. radii and. model magnitudes given by the SDSS catalogs.," We apply a number of, generally very small, corrections to the $\sigma$, radii and model magnitudes given by the SDSS catalogs."
+ The σ measured for earlv-ty galaxies decreases (very slowlv) with distance from the centre.pe and hence the SDSS σ will have a dependeney. on the ratio of the spectrograph aperture to the galaxy’s elective radius.," The $\sigma$ measured for early-type galaxies decreases (very slowly) with distance from the centre, and hence the SDSS $\sigma$ will have a dependency on the ratio of the spectrograph aperture to the galaxy's effective radius."
+ Jorgensen et al. (, rgensen et al. (
+1995) estimated. that mx(Grap£rac)DeL . where ray: isM the de Vaucouleurs model radius in the ραπ.,"1995) estimated that $\sigma\propto (r_{ap}/r_{deV})^{-0.04}$ , where $r_{deV}$ is the de Vaucouleurs model radius in the $r$ -band."
+ As in Bernardi et al. (, As in Bernardi et al. (
+20032) we correct all the σ measurements from the SDSS spectra to apertures of rac/8.,"2003a), we correct all the $\sigma$ measurements from the SDSS spectra to apertures of $r_{deV}/8$."
+" As ro,=1.5 aresec for the SDSS. and racy is in arcsec. this gives 6,5;=οΠο12.0) Od "," As $r_{ap}=1.5$ arcsec for the SDSS, and $r_{deV}$ is in arcsec, this gives $\sigma_{corr}= \sigma_{SDSS}(r_{deV}/12.0)^{-0.04}$ ."
+We correct the de Vaucouleurs mocoel-fit racii. which are semi-major axes. to effective radii (rrr) by multiplving by the square root of the axis ratio. rope=raavb(a.," We correct the de Vaucouleurs model-fit radii, which are semi-major axes, to effective radii $r_{eff}$ ) by multiplying by the square root of the axis ratio, $r_{eff}=r_{deV}\sqrt{b/a}$."
+ Furthermore. there is evidence of svstematic errors in the sky subtraction in the SDSS reductions (e.g. Bernardi et al.," Furthermore, there is evidence of systematic errors in the sky subtraction in the SDSS reductions (e.g. Bernardi et al."
+ 2007: Lauer et al., 2007; Lauer et al.
+ 2007: Hyde Bernardi 2009)., 2007; Hyde Bernardi 2009).
+ These caused the size and total lux of large. extended. objects to be underestimated.," These caused the size and total flux of large, extended objects to be underestimated."
+ By comparing SDSS catalogs and more accurate reductions. Hwde Bernardi (20090). estimated corrections for the SDSS elfective radii.," By comparing SDSS catalogs and more accurate reductions, Hyde Bernardi (2009) estimated corrections for the SDSS effective radii."
+ Phese were positive. a steep function of rope. ancl zero for ρε2.0 arcsec.," These were positive, a steep function of $r_{eff}$, and zero for $r_{eff}\leq2.0$ arcsec."
+ Corrections were also estimated. for. the ce Vaucouleurs model fit magnitudes. these being also a function of rrr. out. zero for rep<1.5 arcsec.," Corrections were also estimated for the de Vaucouleurs model fit magnitudes, these being also a function of $r_{eff}$, but zero for $r_{eff}\leq 1.5$ arcsec."
+" We correct. all our rep, radii as in Hyde Bernareli equation 4.", We correct all our $r_{eff}$ radii as in Hyde Bernardi equation 4.
+ The magnitudes. for these galaxies with lacey=1.0. are the same as de Vaucouleurs moclel-fi magnitudes. except that the radius of the r-band Gt is usec or all passbancds.," The magnitudes, for these galaxies with $\rm frac_{deV}=1.0$ , are the same as de Vaucouleurs model-fit magnitudes, except that the radius of the $r$ -band fit is used for all passbands."
+ Hence we correct our magnitucles in all (ugréz) passbands as in equation 3 of Hyde Bernardi (2005)) but always using the r-band rzp., Hence we correct our magnitudes in all $ugriz$ )passbands as in equation 3 of Hyde Bernardi (2009) but always using the $r$ -band $r_{eff}$.
+ Phis correction wil herefore cause no change in the model-magnitucde colours., This correction will therefore cause no change in the model-magnitude colours.
+ For the magnitudes. noskv-subtraction correction is applied. as these small aperture magnitudes corresponc approximately to a model fit with rr;=1.5 arcsec. where the this correction falls to zero.," For the magnitudes, nosky-subtraction correction is applied, as these small aperture magnitudes correspond approximately to a model fit with $r_{eff}=1.5$ arcsec, where the this correction falls to zero."
+ We calculate. g and. r-bancl A-corrections directly fron the usx-ealibratecl spectra. corrected. for atmospheric and Galactic reddening.," We calculate $g$ and $r$ -band $k$ -corrections directly from the flux-calibrated spectra, corrected for atmospheric and Galactic reddening."
+ To determine the k-correction. for a galaxy at redshift z. we first integrate the galaxy spectrum F(A) over the filter response function.," To determine the k-correction for a galaxy at redshift $z$, we first integrate the galaxy spectrum $F(\lambda)$ over the filter response function."
+ In the g-band. g(A). the apparent magnitude νεος in the AB system. is (loge 2002: Bruzual Charlot 2003) where C(A)=3631(0/A7)11 ere tem 7 is the spectrumof a source with £=3631.0 Jy at all frequencies. defining mig=0.," In the $g$ -band, $g(\lambda)$, the apparent magnitude $g_{spec}$, in the AB system, is (Hogg 2002; Bruzual Charlot 2003) where $C(\lambda)\equiv 3631\, (c/\lambda^2)\times 10^{-23}$ erg $^{-1}$ $^{-2}$ is the spectrumof a source with $F_{\nu}=3631.0$ Jy at all frequencies, defining $m_{AB}=0$."
+ In practice. this integration is performed by summing over the SDSS pixels which are given as £N at intervals of A(logA)=0.0001.," In practice, this integration is performed by summing over the SDSS pixels which are given as $F_{\lambda}$ at intervals of $\Delta({\rm log}~\lambda)=0.0001$."
+ LE we sum over the pixels. then where the summation is over all pixels 9200).," If we sum over the pixels, then where the summation is over all pixels $9200\rm\AA$ )."
+ and we have defined w;=g;/ Mg.Wethen calculate the magnitude go; which would bemeasured in the g-band if the spectrum is de-redshifted bv a factor 1|z. back to z= 0.," and we have defined $w_i \equiv g_i/\Sigma_i g_i$ .Wethen calculate the magnitude $g_{0spec}$ which would bemeasured in the $g$ -band if the spectrum is de-redshifted by a factor $1+z$, back to $z=0$ ."
+" To conserve Lux. this ""compression. increases the Z normalization bv a factor >. making We define ri,! and ros,! similarly. so the A-corrections are"," To conserve flux, this `compression' increases the $F_{\lambda}$ normalization by a factor $1+z$ , making We define $r_{spec}$ and $r_{0spec}$ similarly, so the $k$ -corrections are"
+images (Robertsetal.2001). was resolved into two sources with LIST images (Coadetal.2002).,images \cite{roberts01} was resolved into two sources with HST images \cite{goad02}.
+. For Holmberg LL. Pakull and Mironi (2002) report that the A4686 emission has a EWLIAL of 2.2” and that the Oi] is ollset to the West of the emission.," For Holmberg II, Pakull and Mironi (2002) report that the $\lambda 4686$ emission has a FWHM of $2.2\arcsec$ and that the ] is offset to the West of the emission."
+ Our LIST images confirm these basie results ancl provide much more detailed information on the morphology of the nebula., Our HST images confirm these basic results and provide much more detailed information on the morphology of the nebula.
+" The maximum extent of the emission is to the Wes ol the central star and is about 1.7"" or 26 pe for a distance to Holmberg HE οἱ 3.05 Alpe.", The maximum extent of the emission is to the West of the central star and is about $1.7\arcsec$ or 26 pc for a distance to Holmberg II of 3.05 Mpc.
+ The extent of the nebula towards the East is shorter. about 1.07 or 15 pe.," The extent of the nebula towards the East is shorter, about $1.0\arcsec$ or 15 pc."
+ To the West of the central star (from NW to SW). the morphology of the 1111 emission is similar to. but. more extended: than. tha of the emission and emission is low close to the star and stronger farther out. particularly past. the poin where the emission. peaks.," To the West of the central star (from NW to SW), the morphology of the $\beta$ emission is similar to, but more extended than, that of the emission and ] emission is low close to the star and stronger farther out, particularly past the point where the emission peaks."
+ This is consistent with the behaviour expected for a photo-ionized nebula., This is consistent with the behaviour expected for a photo-ionized nebula.
+ emission is produced. in regions of high excitation ancl should. be concentrated near the central source., emission is produced in regions of high excitation and should be concentrated near the central source.
+ L3 is produced over à broad. range of excitation ancl should cover a larecr spatial range thanHen. Ol], $\beta$ is produced over a broad range of excitation and should cover a larger spatial range than. ]
+ is preferentially produced. at. lower excitation than and should be produced. in the outer regions of the nebula., is preferentially produced at lower excitation than and should be produced in the outer regions of the nebula.
+ Hence. the relative morphologies of theHeu. H3. and Or] emission. support. the hypothesis that the nebula is photoionized. (Pakull&Al," Hence, the relative morphologies of the, $\beta$, and ] emission support the hypothesis that the nebula is photoionized \cite{pakull02}."
+irion, Fig.
+i2002).. Fig. 2 shows the surface brightness versus radius calculated for a photoionized nebula (discussed. further below) for comparison., \ref{cloudy} shows the surface brightness versus radius calculated for a photoionized nebula (discussed further below) for comparison.
+ To the East and South of the central star. the morpholoev. of the Le? emission is very similar to that. of the emission and. ΟΙ emission is absent.," To the East and South of the central star, the morphology of the $\beta$ emission is very similar to that of the emission and ] emission is absent."
+ This is also consistent with a X-rav. photoionized nebula if the column density of the nebula integrated along line of sight outward from the central star. is sullicientIv. small that the entire nebula to the East and South is excited. to produce emission (Pakull&Mirioni2002).," This is also consistent with a X-ray photoionized nebula if the column density of the nebula integrated along line of sight outward from the central star, is sufficiently small that the entire nebula to the East and South is excited to produce emission \cite{pakull02}."
+. We note that there is additional emission to the NW of the central star and beyond the extent of the emission which appears to be part of a ring nebula surrounding a bright star visible in the narrow V band image and unrelated to the nebula., We note that there is additional $\beta$ emission to the NW of the central star and beyond the extent of the emission which appears to be part of a ring nebula surrounding a bright star visible in the narrow V band image and unrelated to the nebula.
+ Also. the isolated bright spot to the SE of the central star appears to be due to incomplete subtraction of a star visible in the narrow V band image.," Also, the isolated bright spot to the SE of the central star appears to be due to incomplete subtraction of a star visible in the narrow V band image."
+ Η the nebula is powered by X-ray. photoionization [rom the central source. then the morphology of the nebula appears inconsistent with narrow beaming of the X-ray emission.," If the nebula is powered by X-ray photoionization from the central source, then the morphology of the nebula appears inconsistent with narrow beaming of the X-ray emission."
+ The nebular emission is likely isotropic., The nebular emission is likely isotropic.
+ Fherefore. the measured [ux should provide a good estimate of the true luminosity.," Therefore, the measured flux should provide a good estimate of the true luminosity."
+ For a distance to Holmboerg LL of 3.05 Alpe. the A4686 line Hux found here for the nebula implies a line luminosity o£ 2.7LO’eres1.," For a distance to Holmberg II of 3.05 Mpc, the $\lambda
+4686$ line flux found here for the nebula implies a line luminosity of $2.7 \times 10^{36} \rm \, erg \, s^{-1}$."
+ Phis is consistent with the value reported by Pakull Mironi (2002)., This is consistent with the value reported by Pakull Mironi (2002).
+ A4686 emission is produced via the recombination of doubly ionized. Le and acts as à photon counter of radiation ionizing the nebula in the — Lyman continuum above 54 eV. Following Pakull Alironi (2002). we use this property to estimate the truce luminosity. of the N-rav source.," $\lambda 4686$ emission is produced via the recombination of doubly ionized $^{++}$ and acts as a photon counter of radiation ionizing the nebula in the $^{+}$ Lyman continuum above 54 eV. Following Pakull Mironi (2002), we use this property to estimate the true luminosity of the X-ray source."
+ The major uncertainty in their estimate was imprecise knowledge of the shape of the photon spectrum illaminating the nebula., The major uncertainty in their estimate was imprecise knowledge of the shape of the photon spectrum illuminating the nebula.
+ Using the NMM-Newton spectra. we are able to reduce this uncertainty.," Using the XMM-Newton spectra, we are able to reduce this uncertainty."
+ We used the photoionization codeCloudy version 94.00 (Ferland2001). to estimate the true X-ray luminosity based on the measured. luminosity., We used the photoionization code version 94.00 \cite{ferland01} to estimate the true X-ray luminosity based on the measured luminosity.
+ We used a metallicity of Z=O0.07Z. (Alirioni2002).," We used a metallicity of $Z = 0.07
+Z_{\sun}$ \cite{mirioni02}."
+. The relation between the total lux and the total ionizing X-ray [lux is not sensitive to the metallicity., The relation between the total flux and the total ionizing X-ray flux is not sensitive to the metallicity.
+ We ran simulations over a range of hydrogen density within the nebula from 1 t0 100 7., We ran simulations over a range of hydrogen density within the nebula from 1 to 100 $^{-3}$.
+" ""Ehe relation between the total fux and the total ionizing X-ray Bux is not sensitive to the density. as long as the region of Lully ionized He is contained within the simulated nebula.", The relation between the total flux and the total ionizing X-ray flux is not sensitive to the density as long as the region of fully ionized He is contained within the simulated nebula.
+ The spatial extent of the emitting region is sensitive to the density., The spatial extent of the emitting region is sensitive to the density.
+ A constant. density of 10 cm gave a reasonable match to the observed. spatial extent of he emission. and we use this value in the simulations »esented here.," A constant density of 10 $^{-3}$ gave a reasonable match to the observed spatial extent of the emission, and we use this value in the simulations presented here."
+ In. the modeling. we assumed. a spherical ecometry with a filling [actor of unity.," In the modeling, we assumed a spherical geometry with a filling factor of unity."
+ We assumed. that here is no absorption between the A-rayv source and the nebula., We assumed that there is no absorption between the X-ray source and the nebula.
+ We also include the photon Hux of an OSV star modelled as a blackbody with a temperature 42000 Ix and a uminosity of 3.2.10eres+.," We also include the photon flux of an O5V star modelled as a blackbody with a temperature 42000 K and a luminosity of $3.2 \times 10^{39} \rm \, erg
+\, s^{-1}$."
+ Fig., Fig.
+ 2 shows the caleulated surface brightness versus apparent olfset from the X-ray source., \ref{cloudy} shows the calculated surface brightness versus apparent offset from the X-ray source.
+ The surface brightness was calculated by integrating he emissivitv along lines of sight through a spherically symmetric nebula., The surface brightness was calculated by integrating the emissivity along lines of sight through a spherically symmetric nebula.
+ ‘Table 1 shows the photoionization luminosities required o produce the measured Luminosity for the various best it Comptonization spectra., Table \ref{specfits} shows the photoionization luminosities required to produce the measured luminosity for the various best fit Comptonization spectra.
+ The inferred. photoionization uminositios range [rom 3.7 to 6.110eres|.," The inferred photoionization luminosities range from $3.7$ to $6.1 \times 10^{39} \rm \, erg \, s^{-1}$."
+ The ow luminosity value comes from the September 2002 observation during which the source was in an unusual ονκο state (Dewanganetal.2004)., The low luminosity value comes from the September 2002 observation during which the source was in an unusual low/soft state \cite{dewangan04}.
+.. The correct. X-ray spectrum to use in the photoionization mocelling would be a luminosity weighted average of observed spectra sampling a duration comparable to the recombination time of in the nebula. roughly 3000 vears for an electron. density of 10 ancl temperature of 20.000 Ix. The Iuminosity weighting would suggest a photoionization luminosity near the high end of the range.," The correct X-ray spectrum to use in the photoionization modelling would be a luminosity weighted average of observed spectra sampling a duration comparable to the recombination time of $^{++}$ in the nebula, roughly 3000 years for an electron density of 10 $^{-3}$ and temperature of 20,000 K. The luminosity weighting would suggest a photoionization luminosity near the high end of the range."
+ ‘To test the sensitivity of these results to the flux from the companion star. we also made runs with the O5V star replaced by a B2 Ib star modelled as a blackbody with a temperature 18.500 Ix. and a luminosity of 2107orgs," To test the sensitivity of these results to the flux from the companion star, we also made runs with the O5V star replaced by a B2 Ib star modelled as a blackbody with a temperature 18,500 K and a luminosity of $2 \times 10^{38} \rm \, erg \, s^{-1}$."
+ The inferred photoionization luminosities increase by 12 to ifa B2 Ib star is used in place of the O5V star., The inferred photoionization luminosities increase by 12 to if a B2 Ib star is used in place of the O5V star.
+ As noted above. our photoionization mocel. assumes a covering factor of unity ancl sullicient depth raclially so that the entire. N-rav αν of the central source is absorbed.," As noted above, our photoionization model assumes a covering factor of unity and sufficient depth radially so that the entire X-ray flux of the central source is absorbed."
+ Inspection of the line emission image shows that the covering factor is below unity., Inspection of the line emission image shows that the covering factor is below unity.
+ Also. as described above. Ou} emission should be present at larger radii than the emission if the nebula is not. density. bounded.," Also, as described above, ] emission should be present at larger radii than the emission if the nebula is not density bounded."
+ Comparison of the O1] anc images shows the nebula is density. bounded: to the East and. South of the central star., Comparison of the ] and images shows the nebula is density bounded to the East and South of the central star.
+ Pherefore. the photoionization Luminosity quoted here should be considered a lower bound to the total luminosity. of the ionizing radiation.," Therefore, the photoionization luminosity quoted here should be considered a lower bound to the total luminosity of the ionizing radiation."
+ Partial absorption intrinsic to the X-ray source would harden the soft. X-ray (ux illuminating 1f nebula ancl would. increase the required Luminosity., Partial absorption intrinsic to the X-ray source would harden the soft X-ray flux illuminating the nebula and would increase the required luminosity.
+ The unabsorbed X-ray luminosity in the Comptonization models. assuming isotropic emission. ranges from 5 to 17 for⋅ the three observations.," The unabsorbed X-ray luminosity in the Comptonization models, assuming isotropic emission, ranges from 5 to $17 \times
+10^{39} \rm \, erg \, s^{-1}$ for the three observations."
+. This Mis consistent. with the luminosity inferred. from the Hell emission. given," This is consistent with the luminosity inferred from the emission, given"
+the GARGS position were rejected and the remainder. were inspected by eve.,the 6dFGS position were rejected and the remainder were inspected by eye.
+ Ehe reader is referred to ο for a detailer description of the radio-optical cross-miatching technique., The reader is referred to \citet{mauch} for a detailed description of the radio-optical cross-matching technique.
+ ? estimated. a reliability of better than 99% anne completeness of better than 96% for NWSS identifications of θα galaxies., \citet{mauch} estimated a reliability of better than $\%$ and completeness of better than $\%$ for NVSS identifications of 6dFGS galaxies.
+" In this context we use the term ""reliability to refer to the number of spurious identifications obtainec (less than 1%) and ‘completeness’ to refer to the number of eeoniune associations müssed (less than 44).", In this context we use the term `reliability' to refer to the number of spurious identifications obtained (less than $\%$ ) and `completeness' to refer to the number of geniune associations missed (less than $\%$ ).
+ Ht is expectec that SUAISS identifications will have similar reliability zu completeness às its resolution at southern declinations is comparable to that of the NVSS., It is expected that SUMSS identifications will have similar reliability and completeness as its resolution at southern declinations is comparable to that of the NVSS.
+ At a740° where the major axis of the SUMSS. bean is TÜ aaresec. NVSS and SUAISS contours were simultaneously overlaid onto optical images ancl inspected. to improve the reliability. of the cross-matching.," At $\delta>-40^\circ$ where the major axis of the SUMSS beam is $>70$ arcsec, NVSS and SUMSS contours were simultaneously overlaid onto optical images and inspected to improve the reliability of the cross-matching."
+ Using this method found that 918 (27%) RASSGdECGS sources are detected in the radio., Using this method found that 918 $27\%$ ) RASS–6dFGS sources are detected in the radio.
+" We define this sample as the RASS6dECGS ""radio sample” and it is discussed in more detail in Section 7.4..", We define this sample as the RASS–6dFGS `radio sample' and it is discussed in more detail in Section \ref{radioprops}.
+ The final data catalogue contains 3405 RASS sources with corresponding optical ancl redshift’ information where available., The final data catalogue contains 3405 RASS sources with corresponding optical and redshift information where available.
+ These are catalogued by their 6dE€GS targetname of the form DDAIAISS*. reflecting the J2000 coordinates.," These are catalogued by their 6dFGS targetname of the form $-$ DDMMSS', reflecting the J2000 coordinates."
+ The content of each column are as (1) ο tareetname: ol the form DDAIAISS. (, The content of each column are as (1) 6dFGS targetname; of the form $-$ DDMMSS. (
+23) Optical position in J2000 coorcinates. (,2–3) Optical position in J2000 coordinates. (
+45) by and Ro magnitudes from the USNO database. (,4–5) $b_{\rm J}$ and R magnitudes from the USNO database. (
+6) RASS name: of the form DDAILAISS. reflecting 2000 coordinates. (,"6) RASS name; of the form $-$ DDMMSS, reflecting J2000 coordinates. ("
+1) Extended flag: “EX means that the X-ray source is extended.,7) Extended flag: `E' means that the X-ray source is extended.
+ This is determined. such that the source extent eiven. in. the RASSBSC --is larger than ⋅∕∕35°. (, This is determined such that the source extent given in the RASS–BSC is larger than $''$. (
+S9) RASS counts per second ancl uncertainty. (,8–9) RASS counts per second and uncertainty. (
+10) Galactic hydrogen column density in units of 7.,10) Galactic hydrogen column density in units of $^{-2}$ .
+ This was calculated using the NLL) program in the HEASARC collection which uses data from? ancl ?.. (, This was calculated using the N(H) program in the HEASARC collection which uses data from \citet{1990ARA&A..28..215D} and \citet{2005A&A...440..775K}. (
+1112) RASS Hux and uncertainty in  +.,11–12) RASS flux and uncertainty in $^{-2}$ $^{-1}$.
+ The tux was calculated: using a fixed. photon index of 1.7 and the Galactic column to determine the unabsorbed [ux in the O.1-24kkeV band., The flux was calculated using a fixed photon index of 1.7 and the Galactic column to determine the unabsorbed flux in the keV band.
+ Vhis made use of the PIAIAIS program in the HIZXASATC collection. (, This made use of the PIMMS program in the HEASARC collection. (
+13) NVSS flux in mJy. (,13) NVSS flux in mJy. (
+14) SUAISS tux in m.s. (,14) SUMSS flux in mJy. (
+1516) Redshift ancl uncertainty from the 6dECS. (,15–16) Redshift and uncertainty from the 6dFGS. (
+17) Quality Hag: “O° signifies that the object wasn't observed as part of the 6.,17) Quality flag; `0' signifies that the object wasn't observed as part of the 6dFGS.
+ Values z3 are regarded as reliable: See Section 3.. (, Values $\geq 3$ are regarded as reliable; See Section \ref{6dfspectra}. (
+18) Programme LD (proglD) of the GdkCGS spectrum.,18) Programme ID (progID) of the 6dFGS spectrum.
+ See Section 7.5., See Section \ref{progid}.
+0 X value of 0. once again signifies that the source wasn't observed. (, A value of `0' once again signifies that the source wasn't observed. (
+19) Classification from NED (20) hedshift from NED (21) Reference of the NED redshitt. (,19) Classification from NED (20) Redshift from NED (21) Reference of the NED redshift. (
+22) Comments - This is à brief note regarding any features in the 6dE€GS spectrum.,22) Comments - This is a brief note regarding any features in the 6dFGS spectrum.
+ See Section 6.1.., See Section \ref{comments}.
+ We have chosen to include the USNO maegnitudes in the catalogue as these were originally used to select. targets prior to the GAGS observations., We have chosen to include the USNO magnitudes in the catalogue as these were originally used to select targets prior to the 6dFGS observations.
+ Optical magnitudes from SuperCOSMOS. have since been included: as part of. the σαCS target sample anc are available from the 6dbCs database (??)..," Optical magnitudes from SuperCOSMOS have since been included as part of the 6dFGS target sample and are available from the 6dFGS database \citep{6dfDR2, 6df2009}."
+ The first 50 entries of the RASSος catalogue are shown in Table 3.., The first 50 entries of the RASS–6dFGS catalogue are shown in Table \ref{examplecat}.
+" Jo prevent these comments. [rom being unneccesarily complicated they are generally self-explanatorv. one-word entries as follows: ""broad! - the spectrum features broad emission lines."," To prevent these comments from being unneccesarily complicated they are generally self-explanatory, one-word entries as follows: `broad' - the spectrum features broad emission lines."
+ ‘narrow - the spectrum has narrow emission features., `narrow' - the spectrum has narrow emission features.
+" ""abs - the spectrum exhibits only absorption lines.", `abs' - the spectrum exhibits only absorption lines.
+" ""BLLac' - a featureless spectrum with strong continuum.", `BLLac' - a featureless spectrum with strong continuum.
+" ""active Mestar - characteristic Al-star spectrum. with strong Balmer emission at z=0.", `active M-star' - characteristic M-star spectrum with strong Balmer emission at z=0.
+" WD? - white dwarf CV - Cataclvsmie variable ""neb! - Spectrum displays >=0 nebula emission ""bad splicing! - refers to an error in the data reduction process when the two arms of the spectrum (blue and red) were not matched together correctly in the final ""(fringing - [ringing occasionally occurs causing strong oscillations in the spectrum.", `WD' - white dwarf `CV' - Cataclysmic variable `neb' - Spectrum displays $z=0$ nebula emission `bad splicing' - refers to an error in the data reduction process when the two arms of the spectrum (blue and red) were not matched together correctly in the final `fringing' - fringing occasionally occurs causing strong oscillations in the spectrum.
+ This is due to either a small air pocket or damage in the fibre whieh causes it to act like a [abry-perot filter., This is due to either a small air pocket or damage in the fibre which causes it to act like a fabry-perot filter.
+ Some redshifts could still be distinguished ‘hlucfrecd arm only - an error with either the red or blue arm during observations resulting in only hall of the spectrum being available., Some redshifts could still be distinguished `blue/red arm only' - an error with either the red or blue arm during observations resulting in only half of the spectrum being available.
+ ‘contamination - a nearby source (generally a foreground star) dominates the spectrum.masking any," `contamination' - a nearby source (generally a foreground star) dominates the spectrum,masking any"
+18.2 yan.,18.2 $\micron$.
+ The lack of any detectable eemission also shows that star formation is weak in the cirenmnuclear region., The lack of any detectable emission also shows that star formation is weak in the circumnuclear region.
+ 4., 4.
+ Assuming that the Iuminositv in NGC 4151 is anisotropic (~13x). the extended mid-IR emission in NGC 4151 is consistent with thermal re-radiation from dust grains in ihe NLR heated by a central engine.," Assuming that the luminosity in NGC 4151 is anisotropic $\thicksim
+13\times $ ), the extended mid-IR emission in NGC 4151 is consistent with thermal re-radiation from dust grains in the NLR heated by a central engine."
+ 5., 5.
+ We place an upper limit on the size of the torus in the mid-IR of <35pe consistent with the measurements of N90. and Raz et.," We place an upper limit on the size of the torus in the mid-IR of $%
+\lesssim 3 5pc consistent with the measurements of N90, and Ruiz et."
+ al. (, al. (
+2002).,2002).
+ This results in an upper limit (o the mid-IR contribution from a dusty torus in NGC 4151 of <73% of the total emission at 10.8 jam and 13.2 jan based on our unresolved (PSF) component.," This results in an upper limit to the mid-IR contribution from a dusty torus in NGC 4151 of $\leq 73\%$ of the total emission at 10.8 $%
+\micron and 18.2 $\micron$ based on our unresolved (PSF) component."
+ 6., 6.
+" Mil-IR. measurements of the proposed torus by N90 as well as upper limits derived from this paper are roughly consistent with the ""onion-skin model of Pedlar et ((1998).", Mid-IR measurements of the proposed torus by N90 as well as upper limits derived from this paper are roughly consistent with the “onion-skin” model of Pedlar et (1998).
+ In this model. ionizing photons in the plane of the torus may be ~10 - 40 times less (han seen from Earth.," In this model, ionizing photons in the plane of the torus may be $\thicksim $ 10 - 40 times less than seen from Earth."
+ We would like to thank the Florida Space Grant Consortium lor funding which led to (he completion of (his work as well as engineer Chris Carter who provided invaluable support while these observations were taken at Gemini North., We would like to thank the Florida Space Grant Consortium for funding which led to the completion of this work as well as engineer Chris Carter who provided invaluable support while these observations were taken at Gemini North.
+scenario mass transfer is driven by the expansion of the Hydrogen burning shell as consequence of the nuclear evolution.,scenario mass transfer is driven by the expansion of the Hydrogen burning shell as consequence of the nuclear evolution.
+" This so-called ""stripped-giant evolutionary model. which is valid for A0.45.. has been applied by King(1993:hereafterK93) o the ease of the BH LMXB V404 Cyg with convincing results."," This so-called 'stripped-giant' evolutionary model, which is valid for $M_c \la 0.45 M_{\odot}$, has been applied by \cite{K93} to the case of the BH LMXB V404 Cyg with convincing results."
+ Using the orbital period of GX 339-4 (P?=1.7557 d) into the K93's equations we obtain with i.=AL.ΔΙ. and me=Alb/AL.., Using the orbital period of GX 339-4 $P=1.7557$ d) into the K93's equations we obtain with $m_c=M_c/M_{\odot}$ and $m_2=M_2/M_{\odot}$.
+" According to K93 Ad, is constrained between where the right hand side condition comes from the Schonberg-Chandrasekhar limit.", According to K93 $M_c$ is constrained between where the right hand side condition comes from the Schonberg-Chandrasekhar limit.
+" we use the condition m,=m» in eq we obtain Al,=MoO.IGGAL..", If we use the condition $m_c=m_2$ in \ref{mc} we obtain $M_c=M_2=0.166M_{\odot}$.
+ This mass is the minimum mass yermitted for the companion and represents the case in which all 1ο Hydrogen shell is burned and the companion becomes a Helium white dwarf., This mass is the minimum mass permitted for the companion and represents the case in which all the Hydrogen shell is burned and the companion becomes a Helium white dwarf.
+ As we show in fig. as the dashed line. this limit case constrains the mass of the BH in |GX339 to Ady>GAL. although us lower limit could be as high as Ay=7.2A/. if we consider ye HOS upper limit to (CAL) (ài. ο. 6.34. ) in the K-correction On the other hand. we obtain Ajo.=LIAL. (AL<O.I8T.M. ) by applying the condition AZ./AJ»~QO.17.," As we show in \ref{all} as the dashed line, this limit case constrains the mass of the BH in GX339 to $M_X \geq 6 M_{\odot}$, although this lower limit could be as high as $M_X \geq 7.2 M_{\odot}$ if we consider the H03 upper limit to $f(M)$ (i. e. $6.3 M_{\odot}$ ) in the K-correction On the other hand, we obtain $M_2\leq1.1M_{\odot}$ $M_c\leq0.187M_{\odot}$ ) by applying the condition $M_c/M_2 \sim 0.17$."
+ This is the maximum mass permitted for the companion through the giant model and sets an upper limit to the mass ratio of q<=0.125 (fig. 5))., This is the maximum mass permitted for the companion through the stripped-giant model and sets an upper limit to the mass ratio of $q\leq0.125$ (fig. \ref{mc}) ).
+ If the companion were as massive ας 1.137. the BH would be more massive than 5.637.., If the companion were as massive as $1.1 M_{\odot}$ the BH would be more massive than $8.6 M_{\odot}$.
+ We list the solutions of the model for GX339-4 in Table I.., We list the solutions of the stripped-giant model for GX339-4 in Table \ref{gx339}.
+" Following K93. the radius. luminosity and mass transfer rate of a stripped-giant Roche lobe filling companion can be calculated as function of the mass of the Helium core according to: For the minimum mass solution we obtain //»=1.5677. and Lo=1.19L. which gives a companion effective temperature Zipp=4831 K. On the other hand. if we consider the maximum mass solution we obtain ο=—2.93/8.. Le=39328584. and Tipp,=4533 K. Although spectroscopic features of the companion star have not been detected yet. SFCOI established a lower limit to the r-band Gunn magnitude of the companion of roc20.4."," Following K93, the radius, luminosity and mass transfer rate of a stripped-giant Roche lobe filling companion can be calculated as function of the mass of the Helium core according to: For the minimum mass solution we obtain $R_2=1.56 R_{\odot}$ and $L_2=1.19L_{\odot}$ which gives a companion effective temperature $T_{eff}=4837$ K. On the other hand, if we consider the maximum mass solution we obtain $R_2=2.93 R_{\odot}$, $L_2=3.25 L_{\odot}$ and $T_{eff}=4533$ K. Although spectroscopic features of the companion star have not been detected yet, SFC01 established a lower limit to the r-band Gunn magnitude of the companion of $r\geq20.4$."
+ We have computed the expected + magnitudes for our derived companions as function of the distance (d) to GX 339-4 which is estimated to be in the range 6<dx15 kpe CHO-D., We have computed the expected $r$ magnitudes for our derived companions as function of the distance $d$ ) to GX 339-4 which is estimated to be in the range $6 \leq d \leq 15$ kpc (H04).
+ For the calculation we have used the mean value for the reddening of hb13~1.2 reported in HO4., For the calculation we have used the mean value for the reddening of $E(B-V)\sim1.2$ reported in H04.
+" According to Schmidt-Kaler(0982). Ry=HR|23PLV) where £2, is the dereddened magnitude which is computed following the r to /? transformation reported in Windhorstetal.(1991)."," According to \cite{SK82}, $R_0=R+2.32 E(B-V)$ where $R_0$ is the dereddened magnitude which is computed following the $r$ to $R$ transformation reported in \cite{Wi91}."
+. For d.~6 kpe we estimated 2].8>r20.6 for the minimum and maximum mass solutions respectively whereas we derive 23.82522.6 if d~15 kpe., For $d\sim 6$ kpc we estimated $21.8 \geq r \geq 20.6$ for the minimum and maximum mass solutions respectively whereas we derive $23.8 \geq r \geq 22.6$ if $d\sim 15$ kpc.
+" Therefore. it is clear that a stripped-giant companion is compatible with the non-detection of the secondary by On the other hand. eg.8 gives a mass transfer rate of 4.9MAL band 7,810IAZ. Hor the minimum and maximum mass solutions respectively."," Therefore, it is clear that a stripped-giant companion is compatible with the non-detection of the secondary by On the other hand, \ref{m2p} gives a mass transfer rate of $4.9\times10^{-11} M_{\odot}$ $^{-1}$ and $7.8 \times10^{-10} M_{\odot}$ $^{-1}$ for the minimum and maximum mass solutions respectively."
+" KKB96 showed that the critical mass transfer rate CAL.) for a system to be a persistent source t1. e. not transient) can be estimated by using the expression where m, is the mass of the compact object in A7. and 37,B the orbital period in units of 3 hours.", KKB96 showed that the critical mass transfer rate $\dot{M_{ct}}$ ) for a system to be a persistent source (i. e. not transient) can be estimated by using the expression where $m_1$ is the mass of the compact object in $M_{\odot}$ and $\frac{P}{3hr}$ the orbital period in units of 3 hours.
+" If we substitute in this equation the orbital period of GX 339-4 and assuming Alx~WAL. we obtain AL;~Ss10""AL. yrὃν, which shows that our values obtained through the stripped-giant model ure consistent with the transient nature of the source."," If we substitute in this equation the orbital period of GX 339-4 and assuming $M_X\sim10M_{\odot}$ we obtain $\dot{M_{ct}}\sim8\times10^{-9} M_{\odot}$ $^{-1}$, which shows that our values obtained through the stripped-giant model are consistent with the transient nature of the source."
+ Homanetal.(2005). detected a maximum X-ray flux during the 2002/2003 outburst of y=3.25.10 “ergs lem ? which results in a X-ray luminosity in the range Ly(peak)~1.4.107 ergs+ depending on the assumed distance., \cite{Ho05} detected a maximum X-ray flux during the 2002/2003 outburst of $F_X=3.25\times10^{-8}$ erg $^{-1}$ $^{-2}$ which results in a X-ray luminosity in the range $L_X(peak) \sim 1.4-8.7\times10^{38}$ erg $^{-1}$ depending on the assumed distance.
+ Applying with 90.1 (Frank.King.&Raine 1992)) we obtain Mp035.L54-10ΛΙ. vr.+., Applying with $\eta\sim0.1$ \citealt{ap}) ) we obtain $\dot{M_1}\sim 0.25-1.54\times10^{-7} M_{\odot}$ $^{-1}$.
+ Therefore. we tind a predicted outburst duty cycle for the outburst in the range 0.32.2.0.10? forthe minimum mass solution and AlsAlO513.2107 for the maximum companion mass case.," Therefore, we find a predicted outburst duty cycle for the outburst in the range $-\dot{M_2}/\dot{M_1} \sim 0.32-2.0\times10^{-3}$ for the minimum mass solution and $-\dot{M_2}/\dot{M_1} \sim 0.51-3.2\times10^{-2}$ for the maximum companion mass case."
+ Classical transient sources te. g. V404 Cyg) only show an outburst every few decades and hence its duty cycles are low (~10. I., Classical transient sources (e. g. V404 Cyg) only show an outburst every few decades and hence its duty cycles are low $\sim10^{-4}$ ).
+ This is clearly not the behaviour showed by GX 339-4 which has undergone four outbursts in the last ~LO years., This is clearly not the behaviour showed by GX 339-4 which has undergone four outbursts in the last $\sim 10$ years.
+ This suggests that Ads is close to Al. for GX 339-4. and therefore a companion mass in the upper part of the proposed mass range seems to better explain the high X-ray activity of this system.," This suggests that $\dot{M_2}$ is close to $\dot{M_{ct}}$ for GX 339-4, and therefore a companion mass in the upper part of the proposed mass range seems to better explain the high X-ray activity of this system."
+ The LMXB XTE 71550-564 was discovered on 1998 by the All Sky Monitor (ASM) onboard the (Smithetal. 1998))., The LMXB XTE J1550-564 was discovered on 1998 by the All Sky Monitor (ASM) onboard the \citealt{Sm98}) ).
+ Oroszetal.(2002:hereafterO02) showed that this system harbours a BH with f(Al)=6.86+O.7LAL. and Ady~ 107..., \cite{Or02} showed that this system harbours a BH with $f(M)=6.86 \pm 0.71 M_{\odot}$ and $M_{X} \sim 10M_{\odot}$ .
+ These authors also classify the companion, These authors also classify the companion
+strong perturbations occurring after every disk passage affect the cluster virial (and its internal kinematics) producing equilibriumdamped oscillations consequentlywhich go on for many dynamical times (Gnedin Ostriker 1999).,strong perturbations occurring after every disk passage affect the cluster virial equilibrium (and consequently its internal kinematics) producing damped oscillations which go on for many dynamical times (Gnedin Ostriker 1999).
+" Finally, the structural evolution of the cluster is also influenced by tides which accelerate the of mass-loss (Gnedin et al."," Finally, the structural evolution of the cluster is also influenced by tides which accelerate the process of mass-loss (Gnedin et al."
+ 1999)., 1999).
+ The overall effect on the processcluster velocity dispersion is therefore extremely complex and not obviously resulting in a heating/freezing., The overall effect on the cluster velocity dispersion is therefore extremely complex and not obviously resulting in a heating/freezing.
+ To evaluate the effect of tides on the velocity dispersion of Pal 14 we ran a set of N-body simulations., To evaluate the effect of tides on the velocity dispersion of Pal 14 we ran a set of N-body simulations.
+ We used the last version of NBSymple (Capuzzo-Dolcetta et al., We used the last version of NBSymple (Capuzzo-Dolcetta et al.
+" 2011), an efficient on a hybrid CPU + GPU platform N-bodyexploiting integratora implementeddouble-parallelization on CPUs and on the hosted Graphic Processing Units (GPUs)."," 2011), an efficient N-body integrator implemented on a hybrid CPU + GPU platform exploiting a double-parallelization on CPUs and on the hosted Graphic Processing Units (GPUs)."
+" The precision is guaranteed resorting to direct summation (to avoid truncation errors in force evaluation), and on the usage of order, time integration methods (Kinoshita et al."," The precision is guaranteed resorting to direct summation (to avoid truncation errors in force evaluation), and on the usage of high order, symplectic time integration methods (Kinoshita et al."
+ high1991; Yoshida symplectic1991)., 1991; Yoshida 1991).
+ In particular the code allows to choose between two different symplectic integrators: a second order algorithm (commonly known as leapfrog) and a much more accurate (but also time consuming) sixth order method., In particular the code allows to choose between two different symplectic integrators: a second order algorithm (commonly known as leapfrog) and a much more accurate (but also time consuming) sixth order method.
+ The effect of the external galactic field is taken into account using an analytical representation of its gravitational potential., The effect of the external galactic field is taken into account using an analytical representation of its gravitational potential.
+ We adopted a leap-frog scheme with a time-step of At=3.7x10yr and a softening length of 0.2 pc (following the prescription of Dehnen Read 2011).," We adopted a leap-frog scheme with a time-step of $\Delta
+t=3.7\times 10^{4}~yr$ and a softening length of 0.2 pc (following the prescription of Dehnen Read 2011)."
+" Such a relatively large time-step and softening length do not affect the accuracy of the simulation, because, as mentioned above, the effects of two-body encounters has been found to be negligible in this cluster even in its innermost region (S11; Beccari et al."," Such a relatively large time-step and softening length do not affect the accuracy of the simulation, because, as mentioned above, the effects of two-body encounters has been found to be negligible in this cluster even in its innermost region (S11; Beccari et al."
+ 2011) and the relaxation times at the, 2011) and the relaxation times at the
+found to decline dramatically after the burst. leading ? to conclude that the maser emitting regions were essentiallv wiped out in a time scale of two weeks by the propagation of the nova shock into the Mira envelope.,"found to decline dramatically after the burst, leading \citet{2011PASJ...63..309D} to conclude that the maser emitting regions were essentially wiped out in a time scale of two weeks by the propagation of the nova shock into the Mira envelope."
+ The RS Oph blast wave produced copious X-rays from at most 3 days after optical cliscovery. when it was first observed. by and.BANG characterized by a steady decline thereafter (2222). ," The RS Oph blast wave produced copious X-rays from at most 3 days after optical discovery, when it was first observed by and, characterized by a steady decline thereafter \citep{2006Natur.442..276S, 2006ApJ...652..629B, 2008ApJ...673.1067N,
+2009ApJ...691..418D}."
+V407 Cyve also developed into an X-ray source. though somewhat more slowly. exhibiting a strong brightening at 20 days or so. à peak near 30 days and a steady decline.," V407 Cyg also developed into an X-ray source, though somewhat more slowly, exhibiting a strong brightening at 20 days or so, a peak near 30 days and a steady decline."
+ ? analyzed the UV. and. X-ray lighteurves together with optical spectra obtained. during 1e first three months of the eruption.," \cite{2011A&A...527A..98S}
+ analyzed the UV and X-ray lightcurves together with optical spectra obtained during the first three months of the eruption."
+ Balmer lines showed steady secular narrowing (seealso?).. and lines of Ca V]. Ve Vll]. Fe XN]. and He IL exhibited asymmetric. proliles attributed to an aspherical expansion. reminiscent of that diagnosed for RS Oph from X-ray lines by 2..," Balmer lines showed steady secular narrowing \citep[see
+also][]{2011MNRAS.410L..52M}, and lines of [Ca V], [Fe VII], [Fe X], and He II exhibited asymmetric profiles attributed to an aspherical expansion reminiscent of that diagnosed for RS Oph from X-ray lines by \citet{2009ApJ...691..418D}."
+ While not. nearly as bright as RS Oph in. N-ravs. V407 (νο was in some ways even more remarkable. with a firm detection bv the Large Area Telescope of variable 5-ràv emission in the 0.1.10 GeV range on 2010 March 10the same day as optical cliscoverythat. persisted. for two weeks (7).," While not nearly as bright as RS Oph in X-rays, V407 Cyg was in some ways even more remarkable, with a firm detection by the Large Area Telescope of variable $\gamma$ -ray emission in the 0.1–10 GeV range on 2010 March 10---the same day as optical discovery—that persisted for two weeks \citep{2010Sci...329..817A}."
+" This was the first ever 5-rav detection of a nova explosion and was attributed by 2. 10 z"" production and. subsequent decay. resulting from. collisions of protons accelerated. in the shock.", This was the first ever $\gamma$ -ray detection of a nova explosion and was attributed by \citet{2010Sci...329..817A} to $\pi^0$ production and subsequent decay resulting from collisions of protons accelerated in the shock.
+ ?. pointed out that short. period svmbiotic novae such as RS Oph are unlikely to. produce s-ravs because of the fast evolution of the blast wave in the higher density environment. closer in to the τος giant. whereas long period svstems such as VL07 Cye provide a much loneer-livecl accelerator.," \citet{2011MNRAS.413L..11L} pointed out that short period symbiotic novae such as RS Oph are unlikely to produce $\gamma$ -rays because of the fast evolution of the blast wave in the higher density environment closer in to the red giant, whereas long period systems such as V407 Cyg provide a much longer-lived accelerator."
+ ? further. noted that V407 Cvg could have been a detectable source of high energy neutrinos., \citet{2010PhRvD..82l3012R} further noted that V407 Cyg could have been a detectable source of high energy neutrinos.
+ The outburst. mechanism for novae is thermonuclear runaway on the WD trggered bv the mass of accreted material exceeding a critical limit (22)..," The outburst mechanism for novae is thermonuclear runaway on the WD triggered by the mass of accreted material exceeding a critical limit \citep{1985ApJ...293L..23S, 1988ApJ...325L..35S}."
+ In two earlier papers we have applied sophisticated hydrodynamic models to the nova explosions on RS Oph (2) and U Sco (2) to provide kev insights into the nature of the explosion and its environment.," In two earlier papers we have applied sophisticated hydrodynamic models to the nova explosions on RS Oph \citep{2009A&A...493.1049O}
+ and U Sco \citep{2010ApJ...720L.195D} to provide key insights into the nature of the explosion and its environment."
+ 1n both sets of mocels. ecireumbinary material proved crucial in explaining the observed: X-ray. emission while providing a degree of collimation to the explosions.," In both sets of models, circumbinary material proved crucial in explaining the observed X-ray emission while providing a degree of collimation to the explosions."
+ Here we describe similar detailed hvdrocdynamic simulations of the V407 (νο explosion. and pay particular attention to the elfects. of circumbinary gas density enhancement ancl the secondary Mira companion on the explosion.," Here we describe similar detailed hydrodynamic simulations of the V407 Cyg explosion, and pay particular attention to the effects of circumbinary gas density enhancement and the secondary Mira companion on the explosion."
+ In Sect., In Sect.
+ 2. we describe the hyerodsnamic mocdel. the numerical setup. and the synthesis of X-rav emission: in Sect.," \ref{s:obsanal} we describe the hydrodynamic model, the numerical setup, and the synthesis of X-ray emission; in Sect."
+ 3. we discuss the results: and finally in Sect.," \ref{s:discuss}
+ we discuss the results; and finally in Sect."
+ 4 we clraw our conclusions., \ref{s:conclusion} we draw our conclusions.
+ The blast wave was modeled. by numerically solving the time-dependent. Ες equations of mass. momentum. and energy conservation. including racliative losses described bv an opticallv-thin plasma ancl thermal conduction: the latter incorporated the ellects of heat [ux saturation.," The blast wave was modeled by numerically solving the time-dependent fluid equations of mass, momentum, and energy conservation, including radiative losses described by an optically-thin plasma and thermal conduction; the latter incorporated the effects of heat flux saturation."
+ Owing to the long timescale of the blast evolution. radiative losses and thermal conduction were found to be more important than in the cases of RS Oph (?7)) and U Sco (7)).," Owing to the long timescale of the blast evolution, radiative losses and thermal conduction were found to be more important than in the cases of RS Oph \citealt{2009A&A...493.1049O}) ) and U Sco \citealt{2010ApJ...720L.195D}) )."
+ The long evolution timescale of the blast. coupled: with large expansion velocities of a few thousand km  rendered the spatial extent too large and computationally demanding to perform an extensive. set of. fully. 3-dimensional. (3D) hvedrocdynamic simulations while still resolving the structure of the immediate binary environment. even with many levels of adaptive mesh refinement.," The long evolution timescale of the blast coupled with large expansion velocities of a few thousand km $^{-1}$ rendered the spatial extent too large and computationally demanding to perform an extensive set of fully 3-dimensional (3D) hydrodynamic simulations while still resolving the structure of the immediate binary environment, even with many levels of adaptive mesh refinement."
+ Ehe computational cost is raised by the inclusion of thermal conduction that is solved explicitly. such a scheme being subject to a rather restrictive stability condition as the thermal conduction timescale is e&cnerallv shorter than the dynamical one (c.g. 2??7)).," The computational cost is raised by the inclusion of thermal conduction that is solved explicitly, such a scheme being subject to a rather restrictive stability condition as the thermal conduction timescale is generally shorter than the dynamical one (e.g. \citealt{2000A&A...362L..41H,
+2005CoPhC.168....1H, 2008ApJ...678..274O, 2010A&A...510A..71O}) )."
+ Given the large computational cost required. by 3D simulations. we adopted the following strategy: we first explored the wide parameter space of the model by adopting a 2-dimensional (2D) cvlindrical. coordinate system (r.z) and. therefore. assuming the svstem to be symmetrical with respect to the axis passing through the WD and the companion star: then. for the set of. parameters found. to best reproduce the observations. we relaxed the hypothesis of axisvmmetry ancl performed a fully 3D. simulation in cartesian ecometrv.," Given the large computational cost required by 3D simulations, we adopted the following strategy: we first explored the wide parameter space of the model by adopting a 2-dimensional (2D) cylindrical coordinate system $(r,z)$ and, therefore, assuming the system to be symmetrical with respect to the axis passing through the WD and the companion star; then, for the set of parameters found to best reproduce the observations, we relaxed the hypothesis of axisymmetry and performed a fully 3D simulation in cartesian geometry."
+ The 2D and 3D models developed here are otherwise similar to the 3D models of ? and ? αμα we refer to those works for further details., The 2D and 3D models developed here are otherwise similar to the 3D models of \citet{2009A&A...493.1049O} and \cite{2010ApJ...720L.195D} and we refer to those works for further details.
+ The caleulations were performed using FLASLL. an adaptive mesh refinement multiphysies code for astrophysical plasmas (2). extended with additional computational modules to handle radiative losses and thermal conduction. (see ?.— for the details of the implementation))," The calculations were performed using FLASH, an adaptive mesh refinement multiphysics code for astrophysical plasmas \citep{Fryxell2000ApJS} extended with additional computational modules to handle radiative losses and thermal conduction (see \citealt{2005A&A...444..505O} for the details of the implementation)."
+ The hycrocynamic equations for compressible eas dynamics are solved using the FLASIIL implementation of the piecewice-parabolie method. (PPM: 7)), The hydrodynamic equations for compressible gas dynamics are solved using the FLASH implementation of the piecewice-parabolic method (PPM; \citealt{1984JCoPh..54..174C}) ).
+ We adopted the system. parameters of 2: these are listed in “Table 1.., We adopted the system parameters of \citet{1990MNRAS.242..653M}; these are listed in Table \ref{t:params}.
+ For the hydrodynamic models. the most important system. parameters are the orbit semi-major axis. the wind mass loss rate and wind terminal velocity.," For the hydrodynamic models, the most important system parameters are the orbit semi-major axis, the wind mass loss rate and wind terminal velocity."
+ ποσο determine the gas density into which the blast occurs. ancl [or a given explosion energy largely. control the subsequent evolution timescale of the resulting shock wave svstem.," These determine the gas density into which the blast occurs, and for a given explosion energy largely control the subsequent evolution timescale of the resulting shock wave system."
+ The mass loss rate and terminal velocity are based on the data of ?.. while the orbital separation was estimated. based. on dust. extinction. changes assumed connected: with orbital modulation.," The mass loss rate and terminal velocity are based on the data of \citet{1985ApJ...292..640K}, while the orbital separation was estimated based on dust extinction changes assumed connected with orbital modulation."
+ ?. estimated orbital separations of 14.0. 15.5 and 16.4 AU for WD masses Ayo=0.5. 1.0 and 1.4 M... respectively: we adopt 15.5 AU and Alp=LOAL..," \citet{1990MNRAS.242..653M} estimated orbital separations of 14.0, 15.5 and 16.4 AU for WD masses $M\rs{WD}=0.5$, 1.0 and 1.4 $M_\odot$, respectively; we adopt 15.5 AU and $M\rs{WD}=1.0\,M_\odot$."
+ We also adopted the ? distance of : kpe. derived. from the ? luminosity ppulsation period relation extrapolated to the W407 Cve period of 745 davs.," We also adopted the \citet{1990MNRAS.242..653M} distance of 2.7 kpc, derived from the \citet{1982MNRAS.199..245G} luminosity pulsation period relation extrapolated to the V407 Cyg period of 745 days."
+ Figure 1. shows an example of the initial conditions adopted in our simulations., Figure \ref{fig1} shows an example of the initial conditions adopted in our simulations.
+ The thermonuclear explosion is initiated by a μαspherical Secdov-type blast wave (7) centered on the WD. with radius ro=LO” km (red circle on the right in the inset panel of Fig. 1)).," The thermonuclear explosion is initiated by a spherical Sedov-type blast wave \citep{1959sdmm.book.....S} centered on the WD, with radius $r_{\rm 0} = 10^8$ km (red circle on the right in the inset panel of Fig. \ref{fig1}) )."
+ In analogy with supernova explosions. the total energy. of the blast. fey. is partitioned so that most of the explosion. energy is kinetic. (in particular we assumed. 1/4 of the energy contained in thermal οποίον and the other 3/4 in kinetic energy as representative: e.g. 7).," In analogy with supernova explosions, the total energy of the blast, $E_{\rm 0}$, is partitioned so that most of the explosion energy is kinetic (in particular we assumed 1/4 of the energy contained in thermal energy and the other 3/4 in kinetic energy as representative; e.g. \citealt{1996ApJ...471..279D}) )."
+ The total mass of the ejecta is AM., The total mass of the ejecta is $M\rs{ej}$.
+ The blast. propagates through the extended outer atmosphere (the wind) of the Mira companion ane is olf-set. from the origin of the wind density. distribution by 15.5 AU (i.e. the system orbital separation: see inset panel of Fig. 1))., The blast propagates through the extended outer atmosphere (the wind) of the Mira companion and is off-set from the origin of the wind density distribution by 15.5 AU (i.e. the system orbital separation; see inset panel of Fig. \ref{fig1}) ).
+ We assumed the gas density in the wind is proportional to 27 (where 2 is the radial distance from, We assumed the gas density in the wind is proportional to $R^{-2}$ (where $R$ is the radial distance from
+ (Iambaryvanetal.2004:Stelzer2006:Robracle&," \citep{Reale2002,Reale2007}. \citep{Hambaryan2004, Stelzer2006, Robrade2009}."
+Schmitt2009).. ~10 105enuo (Pallavicinietal.1995).., $\sim10^{30}$ $10^{33}$ \citep{Pallavicini1995}.
+The carbon content in the envelope of asvinptotic giant branch (AGB) stars is believed lo increase along the spectral sequence A145 3C during this phase of stellar evolution.,The carbon content in the envelope of asymptotic giant branch (AGB) stars is believed to increase along the spectral sequence $\rightarrow$ $\rightarrow$ C during this phase of stellar evolution.
+ The origin of this carbon enhancement is the mixing of He-burning products with matter from the convective envelope through the third dredge-up (EDU) mechanism which can happen after each thermal instability (pulse) of the He-shell (Iben&Renzini1982)., The origin of this carbon enhancement is the mixing of He-burning products with matter from the convective envelope through the third dredge-up (TDU) mechanism which can happen after each thermal instability (pulse) of the He-shell \citep{ibr83}.
+. The recurrence of TDU episodes leads to the creation of a carbon (€) star. defined as an AGB star with a C/O ratio higher than unitv in the envelope.," The recurrence of TDU episodes leads to the creation of a carbon (C) star, defined as an AGB star with a C/O ratio higher than unity in the envelope."
+" Among the C-stars there exists a significant eroup of stars (~15%) named J-tvpe stars (Douigue1954). showing strong ""C-bearing molecule absorptions. which usually implies low ΟΙ ratios (<15) (Lambertοἱal.1986:Abia&Isern.1997:OhnakaTsuji1999)."," Among the C-stars there exists a significant group of stars $\sim 15\%$ ) named J-type stars \citep{bou54} showing strong $^{13}$ C-bearing molecule absorptions, which usually implies low $^{12}$ $^{13}$ C ratios $<15$ ) \citep{lam86,abi97,ohn99}."
+. The location of J-stars in the above AGB spectral sequence is far from clear., The location of J-stars in the above AGB spectral sequence is far from clear.
+ In fact. some authors have located these stars in a dillerent evolutive sequence [rom (hat of the ordinary carbon stars (e.g. Chen Kwok 1993: Lorenz-Martins 1996). or even outside the AGB phase.," In fact, some authors have located these stars in a different evolutive sequence from that of the ordinary carbon stars (e.g. Chen Kwok 1993; Lorenz-Martins 1996), or even outside the AGB phase."
+ J-stars have also been considered as the descendants of the late-R. carbon stars. which have similar spectroscopic characteristics (Llovd-Evans1950).," J-stars have also been considered as the descendants of the late-R carbon stars, which have similar spectroscopic characteristics \citep{llo86}."
+ Theoretically. it is not easy to obtain an AGB star with the chemical peculiarities presented by J-stars.," Theoretically, it is not easy to obtain an AGB star with the chemical peculiarities presented by J-stars."
+ Low ο patios can be obtained in current AGB star models of MZ4 M. if hot H-burning takes place at the bottom of the convective envelope (the so-called hot bottom burning. (Lattanzio1999:Sackmann&Boothrove1992).," Low $^{12}$ $^{13}$ C ratios can be obtained in current AGB star models of $\geq 4$ $_\odot$ if hot H-burning takes place at the bottom of the convective envelope (the so-called hot bottom burning, \citep{lat99,sac92}."
+. ILowever. the performance of the CN-cvcle at the same time destrovs PC and. in consequence. the C/O ratio in the envelope is reduced ancl (he star again becomes O-rich.," However, the performance of the CN-cycle at the same time destroys $^{12}$ C and, in consequence, the C/O ratio in the envelope is reduced and the star again becomes O-rich."
+ Thus. a fime-tuning of (he parameters of the AGB models (mass. mixing-length. mass-loss rate. metallicity. ete.).," Thus, a fine-tuning of the parameters of the AGB models (mass, mixing-length, mass-loss rate, metallicity, etc.),"
+ that determine the chemistry of the envelope. seems {ο be required to obtain a J-star.," that determine the chemistry of the envelope, seems to be required to obtain a J-star."
+ Mixing at the Le-core flash has also been proposed as an alternative scenario to form J-stars., Mixing at the He-core flash has also been proposed as an alternative scenario to form J-stars.
+ In this event an, In this event an
+occurrence of LIATUs with the two cluster models.,occurrence of HMUs with the two cluster models.
+ Assuming that the cluster and the sources are at. fixed. redshifts. this involves two steps: first. for a given cluster one has to integrate over source impact parameters. taking into account finite source size. and second. integrate over the distribution of clusters properties.," Assuming that the cluster and the sources are at fixed redshifts, this involves two steps: first, for a given cluster one has to integrate over source impact parameters, taking into account finite source size, and second, integrate over the distribution of clusters properties."
+" We assume that galaxy. clusters formi a one-parameter family. with the core radius r, of ISC and scale radius rs of NEW being fixed. but. possessing a range of masses."," We assume that galaxy clusters form a one-parameter family, with the core radius $r_c$ of ISC and scale radius $r_s$ of NFW being fixed, but possessing a range of masses."
+ The cluster mass function. dNCAL)fdAl. is. derived. from two observed. relations of X-ray selected clusters.," The cluster mass function, $dN(M)/dM$, is derived from two observed relations of X-ray selected clusters."
+ Dased on a sample of I£MSS clusters. Henry ct al. (," Based on a sample of EMSS clusters, Henry et al. ("
+1992) derive cluster luminosity function to be. dN/dLxxLy. where ay varies with redshift but is approximately 0.3. between the redshifts of 0.15 and. 0.6.,"1992) derive cluster luminosity function to be, $dN/dL_X \propto L_X^{-\alpha_X}$, where $\alpha_X$ varies with redshift but is approximately 0.3 between the redshifts of 0.15 and 0.6."
+" A relation between cluster bolometric Luminosity. derived. from LENXNOSAT cata. and velocity dispersion isgiven by Edec Stewart (1091): Lyx0, 77. "," A relation between cluster bolometric luminosity, derived from EXOSAT data, and velocity dispersion isgiven by Edge Stewart (1991): $L_X \propto \sigma_v^{-2.9}$ ."
+Combining these two relations with the assumption that AZoxστ. we arrive at cluster mass function. The absolute normalization is irrelevant. since we will only be dealing with ratios.," Combining these two relations with the assumption that $M \propto \sigma_v^2$, we arrive at cluster mass function, The absolute normalization is irrelevant, since we will only be dealing with ratios."
+ We need not assume an upper or lower cluster mass limits: the lower mass cutoll is elfectivelv achieved because low mass clusters have a negligible lensing cross-section. while at the upper mass end the lensing cross-section increases slower than the rate at which the numbers of clusters decrease due to the steep slope of their mass function.," We need not assume an upper or lower cluster mass limits: the lower mass cutoff is effectively achieved because low mass clusters have a negligible lensing cross-section, while at the upper mass end the lensing cross-section increases slower than the rate at which the numbers of clusters decrease due to the steep slope of their mass function."
+ We fix the core radius rat tL kpe. which is substantially smaller than the clerivect X-ray ‘core’ sizes. but is consistent with lensing observations (Port Alellicr 1994).," We fix the core radius $r_c$ at 1 kpc, which is substantially smaller than the derived X-ray `core' sizes, but is consistent with lensing observations (Fort Mellier 1994)."
+ For the NEW clusters. we fix ry at 11 kpe. which is close to a typical value obtained in the Navarro et al. (," For the NFW clusters, we fix $r_s$ at 1 kpc, which is close to a typical value obtained in the Navarro et al. ("
+1996) simulations.,1996) simulations.
+ With the physical length scale of the two cluster models fixed. we can now derive the correspondence between the one-parameter ISC and NEW models. ic. we ask what is the relation between αυ and ας of clusters that have the same mass withinps 1 Mpe.," With the physical length scale of the two cluster models fixed, we can now derive the correspondence between the one-parameter ISC and NFW models, i.e. we ask what is the relation between $\kappa_0$ and $\kappa_s$ of clusters that have the same mass within 1 Mpc."
+" EEThe latter is. roughly equal {ο rogo. and. corresponds to 50r. ancl δις, respectively."," The latter is roughly equal to $r_{200}$, and corresponds to $r_c$ and $r_s$, respectively."
+ Using Equations 4.. 5. and 6 we obtain wyz4.6155...," Using Equations \ref{mass_dist_proj1}, \ref{mass_dist_proj2a} and \ref{mass_dist_proj2b} we obtain $\kappa_0\approx 4.615\kappa_s$."
+ We assume that the unlensecd parent population of LIAIUs is the same as that of ares in clusters., We assume that the unlensed parent population of HMUs is the same as that of arcs in clusters.
+ Ehe hall-light radii of the sources of arc images are almost the same as the observed widths of the ares. because the cluster potential is not expected to distort tangential arc images in the racial direction.," The half-light radii of the sources of arc images are almost the same as the observed widths of the arcs, because the cluster potential is not expected to distort tangential arc images in the radial direction."
+ Smail et al. (, Smail et al. (
+1996) measure hall-lieht radii for a sample of δ LIST ares (see their Figure 5).,1996) measure half-light radii for a sample of 8 HST arcs (see their Figure 5).
+ Phe average half- raclius is about 0.5 arcseconds. which is what we will adopt in the present paper.," The average half-light radius is about 0.5 arcseconds, which is what we will adopt in the present paper."
+ We further assume that all the sources are circular with a uniform surface brightness., We further assume that all the sources are circular with a uniform surface brightness.
+ At any given redshift the luminosity function. (LE) of sources is assumed to be a power law. with the slope corresponding to that of the Schechter LE. a.," At any given redshift the luminosity function (LF) of sources is assumed to be a power law, with the slope corresponding to that of the Schechter LF, $\alpha$."
+ The value of à is estimated to be about locally. but may have been steeper in the past. à1.5 (Ellis et al.," The value of $\alpha$ is estimated to be about locally, but may have been steeper in the past, $\alpha\sim 1.5$ (Ellis et al."
+ 1995)., 1995).
+ We use both values below. to account for the possible range of a's depending on the twpe of galaxies and their evolution.," We use both values below, to account for the possible range of $\alpha$ 's depending on the type of galaxies and their evolution."
+ The results are not very sensitive to à., The results are not very sensitive to $\alpha$.
+" We neglect the sources brighter than L,. the characteristic luminosity of Schechter LE. because of their small numbers."," We neglect the sources brighter than $L_\star$, the characteristic luminosity of Schechter LF, because of their small numbers."
+ We need not make any. further assumptions aboutthe source Luminosity function. as we explain below.," We need not make any further assumptions aboutthe source luminosity function, as we explain below."
+" “Phe number of images of type i. where i can be LIAL or ares. for a given cluster characterized byay or ws (or cluster mass AJ). is eiven by. where {4 is the Lleviside step function which is 1 if the image selection criteria are satisfied (Section ??)). and 0 otherwise: Zr, is the faintest observable Luminosity."," The number of images of type i, where i can be HMU or arcs, for a given cluster characterized by$\kappa_0$ or $\kappa_s$ (or cluster mass $M$ ), is given by, where $H_{\rm i}$ is the Heviside step function which is 1 if the image selection criteria are satisfied (Section \ref{criteria}) ), and 0 otherwise; $L_{lim}$ is the faintest observable luminosity."
+ Phe outer integral is over the source impact parameter. y in the source plane.," The outer integral is over the source impact parameter, $y$ in the source plane."
+ Phe integral can be written as The function p(y) is determined by the ISC or NEW model., The integral can be written as The function $\mu(y)$ is determined by the ISC or NFW model.
+ Phe last step in the above equation ds justified because the minimum magnification fr. required. for à detectable image is large (Section ??)). and α is always 71.," The last step in the above equation is justified because the minimum magnification $\mu$ required for a detectable image is large (Section \ref{criteria}) ), and $\alpha$ is always $>$ 1."
+" Both £y7;,, and L« are functions of source recshift. but since we are only interested in the ratio of ngo to nass. the dependency on these quantities cancels out."," Both $L_{lim}$ and $L\star$ are functions of source redshift, but since we are only interested in the ratio of $n_{\rm HMU}$ to $n_{\rm Arcs}$, the dependency on these quantities cancels out."
+ As the results are quite insensitive {ο ος we make no assumptions about the source redshift distribution., As the results are quite insensitive to $z_s$ we make no assumptions about the source redshift distribution.
+ Images are selected based on their size or morphology., Images are selected based on their size or morphology.
+ ‘To be selected as either a giant are or a ΗΔΗ. an image has to be lensecl appreciably.," To be selected as either a giant arc or a HMU, an image has to be lensed appreciably."
+ For an are to be detected. its LAW has to exceed. 10. à commonly used criterium for giant Iuminous ares.," For an arc to be detected, its $L/W$ has to exceed 10, a commonly used criterium for giant luminous arcs."
+ A HMU is defined as an image with central po 10. and LAW« 3.," A HMU is defined as an image with central $\mu>10$ , and $L/W<3$ ."
+ Xdditionallv. the undistorted nature of the image is guaranteed. by requiring that the change in LM. ratio across the image should not exceed 50%..," Additionally, the undistorted nature of the image is guaranteed by requiring that the change in $L/W$ ratio across the image should not exceed ."
+"On the other hand, the long quiescent gap between the two sub-bursts leads us to re-think the 2-stage fallback collapsar scenario that has been used to interpret GRB precursors Mésszárros, 2007).","On the other hand, the long quiescent gap between the two sub-bursts leads us to re-think the 2-stage fallback collapsar scenario that has been used to interpret GRB precursors (Wang Mésszárros, 2007)."
+" In that scenario, the precursor is (Wangproduced by a weak jet formed during the initial core collapse, possibly related to MHD processes associated with a short-lived proto-neutron star, while the main burst is produced by a stronger jet fed by fallback accretion onto the black hole resulting from the collapse of the neutron star."," In that scenario, the precursor is produced by a weak jet formed during the initial core collapse, possibly related to MHD processes associated with a short-lived proto-neutron star, while the main burst is produced by a stronger jet fed by fallback accretion onto the black hole resulting from the collapse of the neutron star."
+" We found that the assumed proto-neutron star rotational energy of a few times 10?! ergs in Wang Mésszárros, 2007 would also be sufficient, when beaming is taken into account, to power the first sub-burst of GRB 110709B. In fact, simple estimates indicate that maximally rotating proto-neutron stars could reach rotational energies as high as several 10°? erg."," We found that the assumed proto-neutron star rotational energy of a few times $10^{51}$ ergs in Wang Mésszárros, 2007 would also be sufficient, when beaming is taken into account, to power the first sub-burst of GRB 110709B. In fact, simple estimates indicate that maximally rotating proto-neutron stars could reach rotational energies as high as several $10^{52}$ erg."
+" Here, we propose a magnetar-to-BH scenario as follows: (1) A magnetar is formed and produces the first sub-burst by releasing its rotation energy via electromagnetic and gravitational radiation in ~10—20 seconds (rest frame)."," Here, we propose a magnetar-to-BH scenario as follows: (1) A magnetar is formed and produces the first sub-burst by releasing its rotation energy via electromagnetic and gravitational radiation in $\sim 10-20$ seconds (rest frame)."
+" A magnetar, rather than a lower field neutron star, is required not only to produce the high luminosity (L+,iso~10°?erg5 1) and (~ 0.6-1 of the first sub-burst (Zhang Mésszárros,Eprest 2001; MeV)Metzger et al."," A magnetar, rather than a lower field neutron star, is required not only to produce the high luminosity $L_{\gamma,iso} \sim 10^{52}\ {\rm erg}\
+{\rm s}^{-1}$ ) and $E_{p,rest}$ $\sim$ 0.6-1 MeV) of the first sub-burst (Zhang Mésszárros, 2001; Metzger et al."
+" 2011), but also to overcome the ram pressure of the"," 2011), but also to overcome the ram pressure of the"
+together with the variability studies presented earlier and observations at other wavelengths im classifvine the NGC 300 sources.,together with the variability studies presented earlier and observations at other wavelengths in classifying the NGC 300 sources.
+ Much of the information is παχΊχος in Table 5.. though some additional explanation ac clarification as regards a few of the individual sources is also given here.," Much of the information is summarized in Table \ref{table_iden}, though some additional explanation and clarification as regards a few of the individual sources is also given here."
+" A point to bear in mined is the subject of backeroun AGN,", A point to bear in mind is the subject of background AGN.
+ An estimate of the number of backeround sources (predominately ACN and QSOs) expected at or above our limiting flux level within a sky area the size of the NGC 300 D25 ellipse (~0.0Lssq. ddeerees) cau be obtained by making use of the logNo5 function from Tasinger al.ss (1998) deep ROSAT ταν survey of tle Lockman Hole., An estimate of the number of background sources (predominately AGN and QSOs) expected at or above our limiting flux level within a sky area the size of the NGC 300 D25 ellipse $\sim0.04$ degrees) can be obtained by making use of the $\log{N}-\log{S}$ function from Hasinger s (1998) deep ROSAT X-ray survey of the Lockman Hole.
+ This uuuber turus out to be about Ὁ. though we note that the nearer their projected positiou to the ceutre of NGC 300. the larder they would be to detect.," This number turns out to be about $-$ 5, though we note that the nearer their projected position to the centre of NGC 300, the harder they would be to detect."
+ One can also compare the surface densitya. of PSPC and URI sources inside and outside of the NGC 300 D25 cllipse. to obtain a perhaps more direct measure of the AGN coutamination within the ealaxy.," One can also compare the surface density of PSPC and HRI sources inside and outside of the NGC 300 D25 ellipse, to obtain a perhaps more direct measure of the AGN contamination within the galaxy."
+. Counting the uunber of unidentified PSPC sources outside of the D25 ellipse aud scaling this for the areas iu question. one obtains a figure of ~9 for the uuuber of unidentified (possibly background. ACN) PSPC. sources expected within a NGC 300 D25 ellipse in this part of the N-rav sky.," Counting the number of unidentified PSPC sources outside of the D25 ellipse and scaling this for the areas in question, one obtains a figure of $\sim$ 9 for the number of unidentified (possibly background AGN) PSPC sources expected within a NGC 300 D25 ellipse in this part of the X-ray sky."
+ The equivalent fleure using the TRI data comes out rather lower. nore like lL.," The equivalent figure using the HRI data comes out rather lower, more like $-$ 4."
+ Problems with this method include the fact that all of he uuideuti&ed non-D25 sources are uulikely to be AGN.4 aud some of course. may be in reality truly associated with NCC 300.," Problems with this method include the fact that all of the unidentified non-D25 sources are unlikely to be AGN, and some of course, may be in reality truly associated with NGC 300."
+ All this said. we expect the ACN contamination within the NCC 300 disk to be anvthine from z [to perhaps more than double this value.," All this said, we expect the AGN contamination within the NGC 300 disk to be anything from $\approx4$ to perhaps more than double this value."
+ The source identification column (€01.55) of Table 5 includes “AGN” cutries. where we are fairly sure that the source is a background. ACN. but also cace.”," The source identification column 5) of Table \ref{table_iden} includes `AGN' entries, where we are fairly sure that the source is a background AGN, but also `acc.'"
+ cutries. where the source appears to be accreting in nature. though whether it is background or related to NGC 300 is difficult {ο sav.," entries, where the source appears to be accreting in nature, though whether it is background or related to NGC 300 is difficult to say."
+ All the sources detected by BPS97 (aud used iu the POO study) are detected., All the sources detected by RPS97 (and used in the P00 study) are detected.
+ All 12 of the newly detecte NGC 300 PSPC sources have a net counts value less than 50 (except for P59. which. though siguificautlv. visible iu the RPS97 study. lav outside thei area of mterest. auk was hence not included in their list).," All 12 of the newly detected NGC 300 PSPC sources have a net counts value less than 50 (except for P59, which, though significantly visible in the RPS97 study, lay outside their area of interest, and was hence not included in their list)."
+ All 11 RPS97 sources re-detected here have a uet counts value ereater than 50. except for sources P29 and P30. the least significaut of the sources listed in RPS97/PO0. (," All 14 RPS97 sources re-detected here have a net counts value greater than 50, except for sources P29 and P30, the least significant of the sources listed in RPS97/P00. ("
+Note RPS97 source 1 corresponds to non-D25 source P21).,Note RPS97 source 14 corresponds to non-D25 source P24).
+ The brightest source. P12. lies within the inner spira armis of NGC 300.," The brightest source, P42, lies within the inner spiral arms of NGC 300."
+ No SNR. eiaut oor ireeions are seen coincident (BLOT). aud it is likely au accreting binary.," No SNR, giant or regions are seen coincident (BL97), and it is likely an accreting binary."
+ Note tha the Eddington limut for a 13. compact object is 1.3<LoS core st. so source DPI2* Nav cinission niv sugeestOO the presceuce of a," Note that the Eddington limit for a $1M_\odot$ compact object is $1.3\times
+10^{38}$ erg $^{-1}$, so source P42's X-ray emission may suggest the prescence of a"
+applied this correction to several nights of data. but it did not provide any improvement to the resulting differential photometry.,"applied this correction to several nights of data, but it did not provide any improvement to the resulting differential photometry."
+ This is likely because we chose reference stars within 5' of the target star over only a part of the detector in which the dome Lats were a good match to the fat-field structure., This is likely because we chose reference stars within $5^{\prime}$ of the target star over only a part of the detector in which the dome flats were a good match to the flat-field structure.
+ As it did not improve the photometry. we did not apply this illumination correction to the data.," As it did not improve the photometry, we did not apply this illumination correction to the data."
+ There is no apparent fringing structure in any of our images which span the wavelength range from 4500-6600A.. including the continuum bands and Ho.," There is no apparent fringing structure in any of our images which span the wavelength range from 4500-6600, including the continuum bands and $\alpha$."
+ Thus. we do not apply a fringe correction.," Thus, we do not apply a fringe correction."
+ In addition. the exposure times ave short and the dark current is negligible.," In addition, the exposure times are short and the dark current is negligible."
+ We do not apply a dark current correction as its application would only add noise to our lighteurves., We do not apply a dark current correction as its application would only add noise to our lightcurves.
+ We obtained. approximately 3300 observations of AU Mic in cach of the continuum filters and ONQO observations in Ho., We obtained approximately 3300 observations of AU Mic in each of the continuum filters and $\sim$ 9800 observations in $\alpha$.
+ After the instrumental signatures were removed. source detection and aperture photometry were performed on all science. frames using the CASU catalogue extraction software. (Irwin&Lewis2001).," After the instrumental signatures were removed, source detection and aperture photometry were performed on all science frames using the CASU catalogue extraction software \citep{IrwinLewis}."
+. The software has been compared with SIExtractor w.ast.cam.ac.uk/wfcam/cdocs/reports/simul/) and. found to be very similar in the completeness. astrometry and photometry tests.," The software has been compared with SExtractor $\sim$ wfcam/docs/reports/simul/) and found to be very similar in the completeness, astrometry and photometry tests."
+ Phe source detection algorithm defines an object as a set of contiguous pixels above a defined detection threshold., The source detection algorithm defines an object as a set of contiguous pixels above a defined detection threshold.
+ Phe routine requires as input the detection limit in units of background σ and the minimum number of connected. pixels above that threshold which define an object., The routine requires as input the detection limit in units of background $\sigma$ and the minimum number of connected pixels above that threshold which define an object.
+ We set a detection threshold of 3 σ and a minimum source size of 15 pixels. so that AU Mic. he brightest star in the field. is casily detected in all the images.," We set a detection threshold of 3 $\sigma$ and a minimum source size of 15 pixels, so that AU Mic, the brightest star in the field, is easily detected in all the images."
+ Potential comparison stars up to 7 magnitudes aunter are also detected., Potential comparison stars up to $\sim 7$ magnitudes fainter are also detected.
+ Aperture photometry with circular apertures was verformed on the detected. objects in all frames., Aperture photometry with circular apertures was performed on the detected objects in all frames.
+ The aperture size allects the precision of the differential wxhotometry. thus the optimum aperture radius was chosen hrough empirical testing of several dillerent sizes.," The aperture size affects the precision of the differential photometry, thus the optimum aperture radius was chosen through empirical testing of several different sizes."
+ An aperture which is too small will be suceptible to centering errors and can be alleeted by the pixelization of the detector. whereas an aperture which is too large will begin to include noise associated with sky pixels. without adding much additional signal.," An aperture which is too small will be suceptible to centering errors and can be affected by the pixelization of the detector, whereas an aperture which is too large will begin to include noise associated with sky pixels without adding much additional signal."
+ Our experience with aperture differential photometry suggests that for bright stars. apertures Larger than the median seeing produce lighteurves with lower ris.," Our experience with aperture differential photometry suggests that for bright stars, apertures larger than the median seeing produce lightcurves with lower rms."
+" The median seeing of our observations is 1.5"" which corresponds to ~2.6 pixels (~ 1.57).", The median seeing of our observations is $\sim 1.5^{\prime\prime}$ which corresponds to $\sim 2.6$ pixels $\sim 1.5^{\prime\prime}$ ).
+ Circular. apertures with radii of 2.8. 4. and δ pixels were tested on a single night of data.," Circular apertures with radii of 2.8, 4, and 8 pixels were tested on a single night of data."
+ The same general features are present in the lighteurves &enerated with all the tested apertures. including the intrinsic stellar. variability. but. the 4. pixel (24/9 aperture produced the AU. Mic lighteurve with the smallest scatter.," The same general features are present in the lightcurves generated with all the tested apertures, including the intrinsic stellar variability, but the 4 pixel $2.4^{\prime\prime}$ ) aperture produced the AU Mic lightcurve with the smallest scatter."
+ Fhus. we obtained instrumental magnitude measurements for all detected: objects on all frames. by summing the background subtracted [lux contained in the 4 pixel aperture.," Thus, we obtained instrumental magnitude measurements for all detected objects on all frames by summing the background subtracted flux contained in the 4 pixel aperture."
+ ‘To derive a differential photometry lighteurve for AU Mic from the instrumental photometry. we first. correct. for atmospheric extinction. using magnitude measurcmicnts obtained on photometric nights for nearby. bright comparison stars.," To derive a differential photometry lightcurve for AU Mic from the instrumental photometry, we first correct for atmospheric extinction using magnitude measurements obtained on photometric nights for nearby, bright comparison stars."
+ In cach filter. we apply a linear least squares fit to the instrumental magnitudes as a function of airmass.," In each filter, we apply a linear least squares fit to the instrumental magnitudes as a function of airmass."
+ Ehe derived. first order extinction. coelIicients are as follows: 0.21 for the 4520/2000) Llter. 0.14 for the 5497/245 filter. anc 0.083. for the 6600/75. and lla filters.," The derived first order extinction coefficients are as follows: 0.21 for the 4520/200 filter, 0.14 for the 5497/245 filter, and 0.083 for the 6600/75 and $\alpha$ filters."
+ We apply these extinction. corrections to all instrumental magnitude measurements., We apply these extinction corrections to all instrumental magnitude measurements.
+ We note that this linear extinction correction is not strictly necessary since it is applied to all stars equally... ancl thus is. removed when the cilferential magnitude is caleulateck.," We note that this linear extinction correction is not strictly necessary since it is applied to all stars equally, and thus is removed when the differential magnitude is calculated."
+ However. we correct For this well understood source of Hux attenuation to derive an instrumental lighteurve which is dominated by the Iluctuating atmospheric transparency. especially as caused bv thin cirrus and other clouds Next. we remove points from the lightcurves which were obtained under very poor. observing conditions (>LO mags below the magnitude obtained under photometric conditions).," However, we correct for this well understood source of flux attenuation to derive an instrumental lightcurve which is dominated by the fluctuating atmospheric transparency, especially as caused by thin cirrus and other clouds Next, we remove points from the lightcurves which were obtained under very poor observing conditions $> 1.0$ mags below the magnitude obtained under photometric conditions)."
+ In acdcition. we remove a handful of measurcments obtained with exposure times «3s which are aclversely alfected by the finite shutter opening time.," In addition, we remove a handful of measurements obtained with exposure times $< 3$ s which are adversely affected by the finite shutter opening time."
+ Finally. we derive the differential photometry lighteurve for AU Mic using. the measurements of ⋅⋅13 comparison⋠ stars withinnac 5/ and up to 5 magnitudes fainter (in the 4520/200 band) han AU Mic.," Finally, we derive the differential photometry lightcurve for AU Mic using the measurements of 13 comparison stars within $\sim 5^{\prime}$ and up to $\sim 5$ magnitudes fainter (in the 4520/200 band) than AU Mic."
+ Then we create asuper comparison star Dighteurve o» averaging the instrumental magnitude liehteurves [or he 13 individual stars (after. subtracting the median value from cach star)., Then we create a comparison star lightcurve by averaging the instrumental magnitude lightcurves for the 13 individual stars (after subtracting the median value from each star).
+ Initially. the 13 reference stars are weighted equally.," Initially, the 13 reference stars are weighted equally."
+ Dillerential photometry liehtcurves are hen obtained for the individual comparison star liehtcurves w subtracting thesuper lighteurve., Differential photometry lightcurves are then obtained for the individual comparison star lightcurves by subtracting the lightcurve.
+ Each comparison star is then assigned a weight value equal to the inverse variance of its dilferential photometry lighteurve normalized so that he weights of all 13 stars sum to one., Each comparison star is then assigned a weight value equal to the inverse variance of its differential photometry lightcurve normalized so that the weights of all 13 stars sum to one.
+ The comparison star lighteurve is then re-created. this time combining the individual lighteurves using a weighted average.," The comparison star lightcurve is then re-created, this time combining the individual lightcurves using a weighted average."
+ The new comparison star lighteurve is then subtracted from the instrumental AU Mic lighteurve., The new comparison star lightcurve is then subtracted from the instrumental AU Mic lightcurve.
+ The resulting cülferential whotometry. lighteurves for AU Mic are shown in Figure 2 for the four filters., The resulting differential photometry lightcurves for AU Mic are shown in Figure \ref{fig:dphot} for the four filters.
+ Since AU. Mic is significantlv brighter han the nearby reference stars used to derive the cüllerential xhotometry. the lighteurve precision is dominated by photon Poisson noise in the reference star magnitude measurements.," Since AU Mic is significantly brighter than the nearby reference stars used to derive the differential photometry, the lightcurve precision is dominated by photon Poisson noise in the reference star magnitude measurements."
+ We have defined an error for cach point on the lighteurve by aking the weighted standard deviation of the 13 comparison star objects used. to derive that point on the AU. Mic ightcurve., We have defined an error for each point on the lightcurve by taking the weighted standard deviation of the 13 comparison star objects used to derive that point on the AU Mic lightcurve.
+ Figure 2. shows the intrinsic stellar variability. previously observed in AW Mic (Iodonoetal.1986)., Figure \ref{fig:dphot} shows the intrinsic stellar variability previously observed in AU Mic \citep{Rodono}.
+. We observe Hare activity and sinusoidal variability indicative of starspots on the stellar surface., We observe flare activity and sinusoidal variability indicative of starspots on the stellar surface.
+ Both sources of variability are [linked to magnetic activity and are. typical for a voung. active M- star. like AU Mic.," Both sources of variability are linked to magnetic activity and are typical for a young, active M-dwarf star, like AU Mic."
+ The instrinsic variability of AU Mic, The instrinsic variability of AU Mic
+caleulation of the DF (Widrow2000).. which in our units is where P?=Mpic! and the parameters are: The results for this svstem are given in Figure 6...,"calculation of the DF \citep{wid00}, which in our units is where $P = \sum_i p_i \varepsilon^i$ and the parameters are: The results for this system are given in Figure \ref{fig:nfw}."
+ As expected. the DII growth induces cusps in both the density and velocity.," As expected, the BH growth induces cusps in both the density and velocity."
+ The logarithmic slope for the density is —7/3. which is (he same as for the self-similar svstem with ὁ=2: thev also both began with the same slope (ox L/H) in the inner region where (the DII disturbance is greatest.," The logarithmic slope for the density is $-7/3$, which is the same as for the self-similar system with $\delta = 2$; they also both began with the same slope $\varrho \propto 1/R$ ) in the inner region where the BH disturbance is greatest."
+ The velocity cusp is (he usual 1//., The velocity cusp is the usual $1/R$.
+" We notice that the velocity tangential anisotropy is also comparable to the svslem with ὁ=2 at small raclii for all black hole masses. aud (hat the velocity perturbation can extend nearly to the core radius for my,>0.1."," We notice that the velocity tangential anisotropy is also comparable to the system with $\delta=2$ at small radii for all black hole masses, and that the velocity perturbation can extend nearly to the core radius for $m_{bh}\ge 0.1$."
+ It is of some interest to explore (he rate at which a black hole woud grow in the above dark matter distributions in view of the black hole mass galactic age relation reported bv Meirilieldefaf(2000)., It is of some interest to explore the rate at which a black hole would grow in the above dark matter distributions in view of the black hole mass galactic age relation reported by \citet{mer00}.
+. Such στον represents an alternative to advection dominated acerelion flow CADAF: see Naravan.Igumenshehev&Abramowiez(2000). ancl references iherein) as à means to grow a black hole invisibly., Such growth represents an alternative to advection dominated accretion flow (ADAF; see \citet{nar00} and references therein) as a means to grow a black hole invisibly.
+ We consider as an illustration the time to erow the black hole bv a factor of ten in the various DFs., We consider as an illustration the time to grow the black hole by a factor of ten in the various DFs.
+" In order to calculate the timescale for the DII (o grow by a [actor of ten. we use the followinge expression for the egrowth of the BIL as it accretes matter: where r, is the radius of the last stable orbit. r,=342,. and Ry,=26M,/c? is the schwarsehild radius of the DII."," In order to calculate the timescale for the BH to grow by a factor of ten, we use the following expression for the growth of the BH as it accretes matter: where $r_x$ is the radius of the last stable orbit, $r_x = 3R_s$, and $R_s = 2 G M_{bh} / c^2$ is the Schwarschild radius of the BH."
+" In our units. and with <7,> given as an integral over the"," In our units, and with $<v_r>$ given as an integral over the"
+"the the last 6 Gyr (z« 0.6), while BH--X suppresses the SFR by 0.8 dex for the last 1.5 Gyr (z« ","the the last 6 Gyr $z<0.6$ ), while BH+X suppresses the SFR by 0.8 dex for the last 1.5 Gyr $z<0.1$ )."
+"BHXRP forms 3.4x10? Mo of stars in the host 0.1).galaxy after z=0.5, and BH+X 3.6x10? Mo, which is less than half of the roughly 7.2x10? Mo formed by the other models."," BHXRP forms $3.4\times10^9$ $_{\odot}$ of stars in the host galaxy after $z=0.5$, and BH+X $3.6\times10^9$ $_{\odot}$, which is less than half of the roughly $7.2\times10^9$ $_{\odot}$ formed by the other models."
+" The BH4-X result is interesting in the context ofL, where the X-ray background produced a strong burst of late star-formation, as gas that had been kept hot through the X-ray background peak at ze2 finally cools and flows to the center of the galaxy to form stars (the so-called “cooling flow"")."," The BH+X result is interesting in the context of, where the X-ray background produced a strong burst of late star-formation, as gas that had been kept hot through the X-ray background peak at $z\approx 2$ finally cools and flows to the center of the galaxy to form stars (the so-called “cooling flow”)."
+" Here, a cooling flow seems to be forming around redshift of 0.3 the lower panel of Figure but the AGN effectively (seeshuts it off."," Here, a cooling flow seems to be forming around redshift of 0.3 (see the lower panel of Figure \ref{fig:sfr-gal-A100}) ), but the AGN effectively shuts it off."
+" The effect with 5)),BHXRP, meanwhile, is clearly related to its enhanced accretion rate after the merger event: we see from Figure 1 that BHXRP accretes less gas during a major merger, leaving a residual which forms a few central stars (see below), but more importantly powers extra feedback for the next several Gyr, suppressing star formation at larger radii."," The effect with BHXRP, meanwhile, is clearly related to its enhanced accretion rate after the merger event: we see from Figure \ref{fig:BHmass} that BHXRP accretes less gas during a major merger, leaving a residual which forms a few central stars (see below), but more importantly powers extra feedback for the next several Gyr, suppressing star formation at larger radii."
+ 'The lower panel of Figure 5 shows the SFR for galaxy A out to a radius of only 3 kpc (the 3D stellar half-mass radius for the galaxy is the range 5—8 kpc)., The lower panel of Figure \ref{fig:sfr-gal-A100} shows the SFR for galaxy A out to a radius of only 3 kpc (the 3D stellar half-mass radius for the galaxy is the range $5-8$ kpc).
+" The results here are somewhat different: while all the models with BHs have no star formation at the present, BHX and BHXRP have nonzero star formation until roughly 1.5 Gyr ago, compared to 3 Gyr for the plain BH model (and 2.5 Gyr for "," The results here are somewhat different: while all the models with BHs have no star formation at the present, BHX and BHXRP have nonzero star formation until roughly 1.5 Gyr ago, compared to 3 Gyr for the plain BH model (and 2.5 Gyr for BH+X)."
+"In terms of mass, BHX forms 105 in the BH4-X).central 3 kpc z= 0.5, 2.8x Mo (the residual from the aftermerger just BHXRPdiscussed), which is 2—3 times more than the 1x105 Mo formed by BH and BH+X, but still less than a tenth of the No BH model (3.1x10? Μο)."," In terms of mass, BHX forms $1.7\times10^8$ $_{\odot}$ in the central 3 kpc after $z=0.5$, BHXRP $2.8\times10^8$ $_{\odot}$ (the residual from the merger just discussed), which is $2-3$ times more than the $1\times10^8$ $_{\odot}$ formed by BH and BH+X, but still less than a tenth of the No BH model $3.1\times10^9$ $_{\odot}$ )."
+" This is an expected result: as inL, when X-ray heating is present, more gas is available at moderate and low redshift (z«2) to be drawn to the galaxy's center in major mergers and in cooling flows, and the AGN feedback, while eventually preventing most of this gas from forming stars, cannot do so immediately."," This is an expected result: as in, when X-ray heating is present, more gas is available at moderate and low redshift $z<2$ ) to be drawn to the galaxy's center in major mergers and in cooling flows, and the AGN feedback, while eventually preventing most of this gas from forming stars, cannot do so immediately."
+" Table 1 shows the stellar mass of the host galaxy for the various models and ICs at three radii: 5 kpc, 30 kpc and 2 Mpc (which is essentially the whole box)."," Table \ref{tab:stars} shows the stellar mass of the host galaxy for the various models and ICs at three radii: 5 kpc, 30 kpc and 2 Mpc (which is essentially the whole box)."
+" The total star formation in the box (column 5) is suppressed slightly (596)) by the presence of the AGN, but the X-rays, whether from the AGN or the background, reduce the total stellar mass by1596, three times as much."," The total star formation in the box (column 5) is suppressed slightly ) by the presence of the AGN, but the X-rays, whether from the AGN or the background, reduce the total stellar mass by, three times as much."
+" However, the background X-rays and, the feedback X-rays cause their suppression differently: as we see in Figure 6,, which shows the star-formation history of the entire box for the halo A runs, the two models with X-ray feedback suppress the star-formation peak at 4>z3, while BH+X is most effective at a somewhat lower redshift, 1.5>z1."," However, the background X-rays and the feedback X-rays cause their suppression differently: as we see in Figure \ref{fig:sfr-big}, which shows the star-formation history of the entire box for the halo A runs, the two models with X-ray feedback suppress the star-formation peak at $4>z>3$, while BH+X is most effective at a somewhat lower redshift, $1.5>z>1$."
+" This clearly reflects the relative timing of the X-ray flux between the feedback and background sources, since the X-ray background peaks in intensity at z&2, while the X-rays from feedback peak when the BH accretion rate does, at 4>z3."," This clearly reflects the relative timing of the X-ray flux between the feedback and background sources, since the X-ray background peaks in intensity at $z\approx2$, while the X-rays from feedback peak when the BH accretion rate does, at $4>z>3$."
+" Thus BHX can have a strong effect on the SFR peak—recall that in Figures 3 4 we saw that BHX has hotter gas than BH+X at z= 3.2—while later when the local AGN is less active, the X-ray background is relatively more effective."," Thus BHX can have a strong effect on the SFR peak—recall that in Figures \ref{fig:gasall}~ \ref{fig:gasvir} we saw that BHX has hotter gas than BH+X at $z=3.2$ —while later when the local AGN is less active, the X-ray background is relatively more effective."
+" This difference in timing also produces a substantial effect on how the stellar mass is distributed,"," This difference in timing also produces a substantial effect on how the stellar mass is distributed,"
+Because uo systematic differences could be deectecl. Ixth night's data we'e combined aud analyzed together (ie.. using the same sensitivity fuuclon and tje samme ext1n«‘1101 law).,"Because no systematic differences could be detected, both night's data were combined and analyzed together (i.e., using the same sensitivity function and the same extinction law)."
+ Bec:use Cilos-100 is a relatively red star (spectral type sdG) tlie variatio between its senstivity calibration aud the othe ‘two bluer stars ixlicated at what waveleeth tle spectra of the due stars were beginning to be allected by second order contariuatloi ( 6860A. altough tlie ati108pjeric “Bo beuxd due to Ov ADSO'ption confuses tlle precise ouset of the elTect).," Because G158-100 is a relatively red star (spectral type sdG) the variation between its sensitivity calibration and the other two bluer stars indicated at what wavelength the spectra of the blue stars were beginning to be affected by second order contamination $\sim$ 6860, although the atmospheric “B” band due to $_2$ absorption confuses the precise onset of the effect)."
+ el lal secoli order contanminatiou is found Wwell efore the AT200 hat oue derives by simply doLet he A3600 vVaveeugth of tle second-order ocking filter., Note that second order contamination is found well before the $\lambda$ 7200 that one derives by simply doubling the $\lambda$ 3600 wavelength of the second-order blocking filter.
+ This is simply because t1e A3600 reler οἱ 1ο poi1it ofthe filter. :uic the profile [9] ‘the trausiulssion o ‘the filter is far TOL1 step-lise.," This is simply because the $\lambda$ 3600 refers to the point of the filter, and the profile of the transmission of the filter is far from step-like."
+ Bevoucl tis wavelength only G158-LOQ was used [or deteruiinine the sensitivity curve (the Πιx of G158-100 is roughly 2 11111itucles weaker a À3500 compared to ATQOO0. so this shoud provicea £&ood calibraticmr out to tje red. unit of our [9]servations at  ATLOO).," Beyond this wavelength only G158-100 was used for determining the sensitivity curve (the flux of G158-100 is roughly 2 magnitudes weaker at $\lambda$ 3500 compared to $\lambda$ 7000, so this should provide a good calibration out to the red limit of our observations at $\sim$ $\lambda$ 7100)."
+ This means ha for enmission lines beween ~ 6860 aand ~ ATLOOA. the fluxes of the emission liies shoulc| uot bea lected. but tha the continuum of the HIE region targets may be contaminated by second order ligtt (thus reducing tlie emission line equivalent wkths).," This means that for emission lines between $\sim$ 6860 and $\sim$ $\lambda$ 7100, the fluxes of the emission lines should not be affected, but that the continuum of the HII region targets may be contaminated by second order light (thus reducing the emission line equivalent widths)."
+ In principle. this is ouly ---uportant [or Using the [He I] A7065 liue to derive a helium abundance if correcious for uuderlyiug absorption are :ipplied ou au equivalent widtli basis.," In principle, this is only important for using the [He I] $\lambda$ 7065 line to derive a helium abundance if corrections for underlying absorption are applied on an equivalent width basis."
+ Because the saudarcl stars were observed over a large range 11 alrinass. during the Ilux calibratiou stage. an extiiction law could be derived f‘om the staricdard star data.," Because the standard stars were observed over a large range in airmass, during the flux calibration stage, an extinction law could be derived from the standard star data."
+ This agreed to within a few percent with he extinction law derived for “the CTIO observaory aud suj»plied within IRAF., This agreed to within a few percent with the extinction law derived for the CTIO observatory and supplied within IRAF.
+ Extractec| spectra are shown for the brightest HII 'egion in each of tle five galaxies in Figures la-e. Eiuissio1 lines were iclentified. aud luxes and e‘rors were measured [followiug the method of Skillinau ]xenuicut (1993: 91999).," Extracted spectra are shown for the brightest HII region in each of the five galaxies in Figures 1a-e. Emission lines were identified, and fluxes and errors were measured following the method of Skillman Kennicutt (1993; SK93)."
+ Values of the logaritlimic exiuctiou at Hj. ο”). wele derived from he error w'eighted average o- deteriuiuatous from the Ha/H2. H+ /He. and Ηδ/Η 3 ratios while siuimltaueously solving for the effects of underlying stelar absorption. EW(HI-abs). (assumed to ye equal in equivalent width for all four Balmer lines).," Values of the logarithmic extinction at $\beta$ , $\beta$ ), were derived from the error weighted average of determinations from the $\alpha$ $\beta$, $\gamma$ $\beta$, and $\delta$ $\beta$ ratios while simultaneously solving for the effects of underlying stellar absorption, EW(HI-abs), (assumed to be equal in equivalent width for all four Balmer lines)."
+" Me assumed the intrinsic case B Baline ‘line ratios calculated by Humuner Storey (LOST). aLd used the redening law of Seaton (1979 as parameterized by Howarth (1983). asstuuine a valtie of R= Ay} /E,32 = 3.2."," We assumed the intrinsic case B Balmer line ratios calculated by Hummer Storey (1987), and used the reddening law of Seaton (1979) as parameterized by Howarth (1983), assuming a value of R $=$ $_{V}$ $_{B-V}$ $=$ 3.2."
+ The assumed reddening law las no associated uucertainty., The assumed reddening law has no associated uncertainty.
+ Uucertaluies 1 CCHS aud were cletermined from Monte Carlo sinulatiois (Olive Skillman 2001)., Uncertainties in $\beta$ ) and EW(HI-abs) were determined from Monte Carlo simulations (Olive Skillman 2001).
+ Figure 2 shows au example solution for one of the observed. HII regiois., Figure 2 shows an example solution for one of the observed HII regions.
+ Note that the errors cerived iL this way are olten siguificaully larger than those derived in the lierature by eitller ASSUME a value for the uucerlyiig absorpionH or derivedH [romH a X> analysH ‘ithout a Monte Carlo analysis of the errors., Note that the errors derived in this way are often significantly larger than those derived in the literature by either assuming a value for the underlying absorption or derived from a $\chi ^2$ analysis without a Monte Carlo analysis of the errors.
+" When tte Baliner ines were corrected for uuderlyi sorption. the hig eritunberec| Baliner lines (H9. H1() HLL. anc H12) were uot assumed to have itical values in terus of ecuivalent widtli but were assiued to luive values of 0.80. 0.65. 0.50. and 0.{0 respectively oftlje lowe: numbered Balmer lines. guided 1wy the low metallicity. instantaneous bi""b models of €onzállez-Delgado. Leitherer. Heckinan (1999)."," When the Balmer lines were corrected for underlying absorption, the higher numbered Balmer lines (H9, H10, H11, and H12) were not assumed to have identical values in terms of equivalent width, but were assumed to have values of 0.80, 0.65, 0.50, and 0.40 respectively of the lower numbered Balmer lines, guided by the low metallicity, instantaneous burst models of Gonzállez-Delgado, Leitherer, Heckman (1999)."
+ Note that all ofthe Sculptor Croup dls lie at Galactic latit«les more negative than —69. aud thus. Galactic extinction along the lines of sight to the H IL τος]us should be ueplieible (e.g.. the Galactic foregrouud E(B-V) for the Sculptor dwarf irregular galxy. Which lies uear the midcle of the projecteddistribution. is caleulated to be 0.012 mag by Seilegel et 11995).," Note that all of the Sculptor Group dIs lie at Galactic latitudes more negative than $-69$, and thus, Galactic extinction along the lines of sight to the H II regions should be negligible (e.g., the Galactic foreground E(B-V) for the Sculptor dwarf irregular galaxy, which lies near the middle of the projecteddistribution, is calculated to be 0.012 mag by Schlegel et 1998)."
+"mass ""unction (as we do) to determine the effect of non-Catssianitv.",mass function (as we do) to determine the effect of non-Gaussianity.
+ Thus. we can directly compare these effects witi the results of non-equipartition. free of biases iutrocdiwed by the choice of mass function. ixd estimate the bias non-equipartition will introduce iu (fforts to detect ton-Gaussian fluctuations with clusters.," Thus, we can directly compare these effects with the results of non-equipartition, free of biases introduced by the choice of mass function, and estimate the bias non-equipartition will introduce in efforts to detect non-Gaussian fluctuations with clusters."
+ Tf «ds fitted :wea constant parameter (ay frozen at zero). then depeiding on the calibration methods. nou-equipartitiou effect can introduce a few percent biases ou the measured cosmological parameters £).," If $w$ is fitted as a constant parameter $w_1$ frozen at zero), then depending on the calibration methods, non-equipartition effect can introduce a few percent biases on the measured cosmological parameters )."
+" Either self-calibratiug ie Y M relation using a power-law form 2) or caibrating the 3 AL relation correctly at low redshift 3) can significant reduce the biases In O3, Ty. Or t» tosl**X."," Either self-calibrating the $Y$ $M$ relation using a power-law form ) or calibrating the $Y$ $M$ relation correctly at low redshift ) can significant reduce the biases in $\Omega_M$, $\sigma_8$, or $w_0$ to $\lesssim 1\%$."
+Ifswe allow c to vary with redshift. the norequipartition effect can introduce a bias in cosinologica paralcter of up to ~105€ DD.," If we allow $w$ to vary with redshift, the non-equipartition effect can introduce a bias in cosmological parameter of up to $\sim 10\%$ )."
+ Iu particular. non-equipartition effects ca1 lutroduce an apparent evolujon iu « of a few percent iu all of th Cases we considered.," In particular, non-equipartition effects can introduce an apparent evolution in $w$ of a few percent in all of the cases we considered."
+ Using the caster abundance techliuicue aloue. an N-ray survey with Sh Nav clusters has already constrained ποne cosmological parameCrs down Τε level in satistical uucertaüntv (Vikhliniuotal. 20001.," Using the cluster abundance technique alone, an X-ray survey with 85 X-ray clusters has already constrained some cosmological parameters down to level in statistical uncertainty \citep{Vik+09}."
+.. Ougoing aud uture SZ surveys will detect thousands of clusters (c.g.Birkinshaw1999:Carl-strometal.2002:Bartlett2 s).. and this will sjeuificauflv muiprove 1ο constraints on cosmological paralcters.," Ongoing and future SZ surveys will detect thousands of clusters \citep[e.g.,][]{Bir99,CHR02,BCM+08}, and this will significantly improve the constraints on cosmological parameters."
+ Therefore. it is muportant to control the systematic uncertainties of galaxy chster plysics at even a percelnage leve].," Therefore, it is important to control the systematic uncertainties of galaxy cluster physics at even a percentage level."
+ Iverodvuamic simulations assuniug equipartition suggest fiat the integrated Y ds a robust nass proxy even when galaxy clusters ave in the process of iereig.oO and lence he iucerated Y is taken to be a rearly 1eal probe for c“OSLOogical studies (Poole 2007: W SR).," Hydrodynamic simulations assuming equipartition suggest that the integrated $Y$ is a robust mass proxy even when galaxy clusters are in the process of merging, and hence the integrated $Y$ is taken to be a nearly ideal probe for cosmological studies \citealt{PBM+07}; \citetalias{WSR+08}) )."
+ Qur results show tha if the non-cquipartition effect is rot properly taken iuto accouit. cosmological parameters can be djased significantly (up to ~ 105).," Our results show that if the non-equipartition effect is not properly taken into account, cosmological parameters can be biased significantly (up to $\sim 10\%$ )."
+ Iu order o fake hne non-equipartitioi effect iuto. account when using cluster abundance to study precision cosimiologx. he ultimate solution is to inclide the non-equipartition effect imm cosmiological sinulatious assuuiug the non-equipartition plysics is known accurately.," In order to take the non-equipartition effect into account when using cluster abundance to study precision cosmology, the ultimate solution is to include the non-equipartition effect in cosmological simulations assuming the non-equipartition physics is known accurately."
+ Tf higher resolution is needed. another approach is to correct he non-cquipartition effect dy performing idealized siauulatious σοιWong&Saraziu2009) or to re-simulate represeutative clusters taken from cosinological sinulatious inchiding the non-equipartition effect with realistic assumptions.," If higher resolution is needed, another approach is to correct the non-equipartition effect by performing idealized simulations \citep[e.g.,][]{WS09} or to re-simulate representative clusters taken from cosmological simulations including the non-equipartition effect with realistic assumptions."
+ For the latter case. the non-equipartition effect can be taken into account together| with other physical processes (e... gas depletion processes during the formation) which may also affect the Y versus ass relation.," For the latter case, the non-equipartition effect can be taken into account together with other physical processes (e.g., gas depletion processes during the formation) which may also affect the $Y$ versus mass relation."
+" Πωπονοα, the above calibration methods by nuuerical simulations rely ou the assumption that the non-equipartition plivsies is known accurately. wuch is in act not the case at present."," However, the above calibration methods by numerical simulations rely on the assumption that the non-equipartition physics is known accurately, which is in fact not the case at present."
+ One of the key systematic uncertainties is the electron heating efficiency at the collisionless shock. j.," One of the key systematic uncertainties is the electron heating efficiency at the collisionless shock, $\beta$."
+ One wav to constraln hne nonm-eqiIpartition plysics is to make direct observalous of accretion shocks. which is currently not feasible.," One way to constrain the non-equipartition physics is to make direct observations of accretion shocks, which is currently not feasible."
+ We may cousrain non-equipartition plysics based on observations of other astroplivsical shocks sucht ax nerecrs shocks and supernova reninauts;, We may constrain non-equipartition physics based on observations of other astrophysical shocks such as mergers shocks and supernova remnants.
+ Iowever. we have to assume these results apply to cluster accretion shocks. which may or may not be the case.," However, we have to assume these results apply to cluster accretion shocks, which may or may not be the case."
+ Another route iuieht be to perforii plasiua. simulations (6.9.. particle-iu-cell «inulatious) to constrain the shock physics.," Another route might be to perform plasma simulations (e.g., particle-in-cell simulations) to constrain the shock physics."
+ However. to apply the plasma simuation results to cluster accretion shocks. a detailed kuowledge of the xe-hock. plivsies such as the magnetic field structure uieht be needed.," However, to apply the plasma simulation results to cluster accretion shocks, a detailed knowledge of the pre-shock physics such as the magnetic field structure might be needed."
+ Clearly. all of the avove calculations are Lecessary to determine the rauge of the svstcmatic uncertainties aud the effects of the lol-equipartition oivsies. aud Τε»coustrain the form ofthe Y versus mass relation.," Clearly, all of the above calculations are necessary to determine the range of the systematic uncertainties and the effects of the non-equipartition physics, and to constrain the form ofthe $Y$ versus mass relation."
+ These should be studied iu he near future., These should be studied in the near future.
+" Until maiuerical suuations can directly determine the effects of non-equipartition on the Y versus mass relation from first principles. we suggest either to selfcalibrate the Y versus inass relation using a power-law form at cach redshift biu 2), or to calibrae the Y versus mass relation corre‘tly at low redshift 3)."," Until numerical simulations can directly determine the effects of non-equipartition on the $Y$ versus mass relation from first principles, we suggest either to self-calibrate the $Y$ versus mass relation using a power-law form at each redshift bin ), or to calibrate the $Y$ versus mass relation correctly at low redshift )."
+" These will reduce the biases due to nou-equipartition on 3.σα, ΟΥ wy to better thau"," These will reduce the biases due to non-equipartition on $\Omega_M, \sigma_8$, or $w_0$ to better than."
+ huportaut biases introduced by other plysical Xocesses can be corrected im addition to the non-equip:wtition correction., Important biases introduced by other physical processes can be corrected in addition to the non-equipartition correction.
+ However. if one would tiv to coustrali the evolution in « to better than together with tιο self-calibration method. the constraints from umiuerica simulations with uncertainties less thia a percent level iüght be necessary.," However, if one would try to constrain the evolution in $w$ to better than, together with the self-calibration method, the constraints from numerical simulations with uncertainties less than a percent level might be necessary."
+ We have shewn that using the cluster abundance to constraint the dark οποίον equation of state requires a full ouxcrstauding of the systematic uncertainties in ealaxy cluster pliysics., We have shown that using the cluster abundance to constraint the dark energy equation of state requires a full understanding of the systematic uncertainties in galaxy cluster physics.
+" Even though we are ouly considering thiο systematic uncertainties introduced by the non-ecuipartition effect. our results also suggest that svstenatic uucertaitics in the Y A relation introduced by other plivsies of even a few percent can introduce a ος1uparable level of biases iu cosmological parameter nie:wurenments,"," Even though we are only considering the systematic uncertainties introduced by the non-equipartition effect, our results also suggest that systematic uncertainties in the $Y$ $M$ relation introduced by other physics of even a few percent can introduce a comparable level of biases in cosmological parameter measurements."
+" Future cluster surveys aiming to constrain departures from ecneral relativity will need to control svsCluatic uncertaimties down to a sub-perceutaec| level (Schuuidtotal.2009).. ancl hence cluster plysics nist be understood iu a conrarable accuracy,"," Future cluster surveys aiming to constrain departures from general relativity will need to control systematic uncertainties down to a sub-percentage level \citep{SVH09}, and hence cluster physics must be understood in a comparable accuracy."
+ Future theoretical calculations aud αποΊσα simulations slild pay particular attention to the effects of non-thermal physies ou the electron pressure profiles., Future theoretical calculations and numerical simulations should pay particular attention to the effects of non-thermal physics on the electron pressure profiles.
+ Poteutial svstematic uncertainties include coreuction. turbulent pressure. magnetic pressure. and relativistic pressure supported bv cosmic rays.," Potential systematic uncertainties include conduction, turbulent pressure, magnetic pressure, and relativistic pressure supported by cosmic rays."
+ Deep observations should also be carried out to coustrain all these effects iu detail for iudividual clusters., Deep observations should also be carried out to constrain all these effects in detail for individual clusters.
+ The outer regions of galaxy clusters are ideal sites for study uou-thermal plysics., The outer regions of galaxy clusters are ideal sites for study non-thermal physics.
+ These studies not only cau increase our undorstaudius of cosimologv. but also can provide information on the plivsics of gaaxy clusters and plasimua plysics under extreme condilonis.," These studies not only can increase our understanding of cosmology, but also can provide information on the physics of galaxy clusters and plasma physics under extreme conditions."
+ KW. thanss Avi Loeb and Briau Mason for helpful discussious., K.W. thanks Avi Loeb and Brian Mason for helpful discussions.
+ Support for this work was provided bv the National Acronautics and Space Acuuinistratio1. through Chandra Award Numbers TMT- CO9-H35N. and GO9-OLISN. NASA NATL ας NNNOSAZ3 and NNNOSAWRAC. aud NASA Suzaku αςLG NNNOSAZOOG. NNNOQATIZ5G. aud NNXOOATITLC. We thank the referee for Lehful comments.," Support for this work was provided by the National Aeronautics and Space Administration, through Award Numbers TM7-8010X, GO9-0135X, and GO9-0148X, NASA XMM-Newton Grants NNX08AZ34G and NNX08AW83G, and NASA Suzaku Grants NNX08AZ99G, NNX09AH25G, and NNX09AH74G. We thank the referee for helpful comments."
+of the statistical test used. when considering only the highest energy (>10? eV at Ia) the correlation between UHECRs and QSOs is consistent with a random distribution at the 1c level.,"of the statistical test used, when considering only the highest energy $> 10^{19.9}$ eV at $\sigma$ ) the correlation between UHECRs and QSOs is consistent with a random distribution at the $1\sigma$ level."
+ In a series of recent papers. Tinyakov and Tkachev (2001. 2002. 2003) claim a correlation between the arrival directions of UHECRs and BL Laes. a subgroup of the QSO sample previously considered.," In a series of recent papers, Tinyakov and Tkachev (2001, 2002, 2003) claim a correlation between the arrival directions of UHECRs and BL Lacs, a subgroup of the QSO sample previously considered."
+ Specifically. the BL Lacs chosen were those identified in the (9th-Edition) Veron-Cetty and Veron (2000) catalogue of Quasars and Active Galactic Nuclei. with redshift z>0.1 or unknown. magnitude η<18. and radio flux at 6 GHz Fy>0.17 Only 22 objects fulfill such restrictions.," Specifically, the BL Lacs chosen were those identified in the (9th-Edition) Veron-Cetty and Veron (2000) catalogue of Quasars and Active Galactic Nuclei, with redshift $z > 0.1$ or unknown, magnitude $m < 18$, and radio flux at 6 GHz $F_6 >
+0.17$ Only 22 objects fulfill such restrictions."
+ In this analysis there is no buffer against contamination by mismeasured protons piled up at the GZK energy limit., In this analysis there is no buffer against contamination by mismeasured protons piled up at the GZK energy limit.
+ The CR sample of Tinyakov and Tkachev consists of 26 events measured by the Yakutsk experiment with energy >10/95 ev (Afanasiev et al., The CR sample of Tinyakov and Tkachev consists of 26 events measured by the Yakutsk experiment with energy $> 10^{19.38}$ eV (Afanasiev et al.
+ 1996). and 39 events measured by the AGASA experiment with energy >10!°°* eV (Hayashida et al.," 1996), and 39 events measured by the AGASA experiment with energy $> 10^{19.68}$ eV (Hayashida et al."
+ 2000)., 2000).
+" The evidence supporting their claim is based on 6 events reported by the AGASA Collaboration (all with average energy <[01 eV). and 2 events recorded with the Yakutsk experiment (both with average energy <10'°° eV). which were found to be within 2.5"" of 5 BL Lacs contained in the restricted sample of 22 sources."," The evidence supporting their claim is based on 6 events reported by the AGASA Collaboration (all with average energy $< 10^{19.9}$ eV), and 2 events recorded with the Yakutsk experiment (both with average energy $< 10^{19.6}$ eV), which were found to be within $2.5^\circ$ of 5 BL Lacs contained in the restricted sample of 22 sources."
+ The chance probability for this coincidence set-up ts found to be 2«107. One drawback of the claim made by Tinyakov and Tkachev (2001) is that the data set used to make the initial assertion is also being used in the hypothesis testing phase., The chance probability for this coincidence set-up is found to be $2 \times 10^{-5}.$ One drawback of the claim made by Tinyakov and Tkachev (2001) is that the data set used to make the initial assertion is also being used in the hypothesis testing phase.
+ Note that if enough searches are performed on a finite data set which is sampled from an isotropic distribution. some highly significant positive results are certain to occur due to the statistical fluctuations that necessarily arise in any finite sampling.," Note that if enough searches are performed on a finite data set which is sampled from an isotropic distribution, some highly significant positive results are certain to occur due to the statistical fluctuations that necessarily arise in any finite sampling."
+ Evans. Ferrer and Sarkar (2002) already called into question whether the selection criteria for the subset of brightest BL Lacs are unbiased.," Evans, Ferrer and Sarkar (2002) already called into question whether the selection criteria for the subset of brightest BL Lacs are unbiased."
+ Strictly speaking. Tinyakov and Tkachev imposed arbitrary cuts on the BL Lae catalogue so as to maximize the signal-to-noise ratio. compensating the different cut adjustments by inclusion of a penalty factor.," Strictly speaking, Tinyakov and Tkachev imposed arbitrary cuts on the BL Lac catalogue so as to maximize the signal-to-noise ratio, compensating the different cut adjustments by inclusion of a penalty factor."
+ Without these arbitrary cuts. the significance of the correlation signal is reduced at the Io level (Evans. Ferrer Sarkar 2002).," Without these arbitrary cuts, the significance of the correlation signal is reduced at the $\sigma$ level (Evans, Ferrer Sarkar 2002)."
+ Moreover. even in acceptance of thisposteriori approach. the estimated value of the penalty factor is subject to debate (Evans. Ferrer Sarkar 2002; Tinyakov Tkachev 2003).," Moreover, even in acceptance of this approach, the estimated value of the penalty factor is subject to debate (Evans, Ferrer Sarkar 2002; Tinyakov Tkachev 2003)."
+ Given the pivotal role played by the penalty factor in testing the hypothesis with a single set of data. it is of interest to circumvent this ambiguity by performing a blind analysis.," Given the pivotal role played by the penalty factor in testing the hypothesis with a single set of data, it is of interest to circumvent this ambiguity by performing a blind analysis."
+ We have at our disposal the cosmic ray arrival directions of the Haverah Park (Stanev et al., We have at our disposal the cosmic ray arrival directions of the Haverah Park (Stanev et al.
+ 1995) and Voleano Ranch (Linsley 1980) experiments. which. although not useful to distinguish a positive correlation (because the penalties involved are probably already as large as the signal which one expects to test). they provide the framework to disregard the correlation 1f none is found in the data.," 1995) and Volcano Ranch (Linsley 1980) experiments, which, although not useful to distinguish a positive correlation (because the penalties involved are probably already as large as the signal which one expects to test), they provide the framework to disregard the correlation if none is found in the data."
+ Surface arrays in stable operation have nearly continuous observation over the entire year. yielding a uniform exposure in right ascension.," Surface arrays in stable operation have nearly continuous observation over the entire year, yielding a uniform exposure in right ascension."
+ However. the declination distribution is different for each experiment. because the relative efficiency of the detection of events depends upon the latitude of the array and detector type.," However, the declination distribution is different for each experiment, because the relative efficiency of the detection of events depends upon the latitude of the array and detector type."
+ As shown by Uchihort et al. (, As shown by Uchihori et al. (
+2000). the field of view of AGASA + Yakutsk Is roughly equal to that of Volcano Ranch + Haverah Park.,"2000), the field of view of AGASA + Yakutsk is roughly equal to that of Volcano Ranch + Haverah Park."
+ It is noteworthy that even though the energy of the Haverah Park events has been reduced by about (Ave et al., It is noteworthy that even though the energy of the Haverah Park events has been reduced by about (Ave et al.
+ 2003). the 27 events contained in our sample. originally with energy >10/75 eV (Lawrence. Reid Watson 1991). are well above the energy cut for Yakutsk’s events selected by Tinyakov and Tkachev.," 2003), the 27 events contained in our sample, originally with energy $> 10^{19.6}$ eV (Lawrence, Reid Watson 1991), are well above the energy cut for Yakutsk's events selected by Tinyakov and Tkachev."
+ Combined with the 6 events recorded at the Volcano Ranch with energy >10139 eV (Linsley 1980). we have a virgin data-set of 33 events. amounting to half of the cosmic-ray arrival directions used to make the claim.," Combined with the 6 events recorded at the Volcano Ranch with energy $> 10^{19.6}$ eV (Linsley 1980), we have a virgin data-set of 33 events, amounting to half of the cosmic-ray arrival directions used to make the claim."
+ In Fig., In Fig.
+ 1. we plot the position on the sky in galactic coordinates of both the UHECRs and the selected BL Lacs., \ref{skymap} we plot the position on the sky in galactic coordinates of both the UHECRs and the selected BL Lacs.
+ There are no positional coincidences between these two samples up to an angular bin —57., There are no positional coincidences between these two samples up to an angular bin $> 5^\circ$.
+ Such an angular scale is well beyond the error in arrival determination. which is found to be z3° (Uchihori et al.," Such an angular scale is well beyond the error in arrival determination, which is found to be $\approx
+3^\circ$ (Uchihori et al."
+ 2000)., 2000).
+ On the basis of the strongly correlated sample analyzed by Tinyakov and Tkachev. one expects the distribution describing the correlation between the set of BL Laces and any UHECR data-set with 33 entries to be Poisson with mean z4.06. Taking the data at face value. this implies a 260 deviation effect.," On the basis of the strongly correlated sample analyzed by Tinyakov and Tkachev, one expects the distribution describing the correlation between the set of BL Lacs and any UHECR data-set with 33 entries to be Poisson with mean $\approx 4.06.$ Taking the data at face value, this implies a $2\sigma$ deviation effect."
+ Moreover. the CL interval of the distribution which samples the correlation between the BL Lacs and cosmic rays recorded by Voleano Ranch + Haverah Park is (0. 3.09) (see. e.g. Feldman Cousins 1998).," Moreover, the CL interval of the distribution which samples the correlation between the BL Lacs and cosmic rays recorded by Volcano Ranch + Haverah Park is (0, 3.09) (see, e.g. Feldman Cousins 1998)."
+ Therefore. the probability to measure the expected mean value z4.06 Is 5656.," Therefore, the probability to measure the expected mean value $\approx 4.06$ is $\ll 5\%$ ."
+ Allin all. the 8 coincidences in the Tinyakov and Tkachev (2001) analysis do not represent a statistically significant effect.," All in all, the 8 coincidences in the Tinyakov and Tkachev (2001) analysis do not represent a statistically significant effect."
+ On a similar track. Gorbunov et al. (," On a similar track, Gorbunov et al. ("
+2002) claimed that a set of 3-ray loud BL Lac objects can be selected by intersecting the EGRET and BL Laes catalogs.,2002) claimed that a set of $\gamma$ -ray loud BL Lac objects can be selected by intersecting the EGRET and BL Lacs catalogs.
+ The only requirement that Gorbunov et al., The only requirement that Gorbunov et al.
+ considered for a BL Lae to be physically associated with an EGRET source is that the angular distance between the best estimated position. of the pair does not exceed 2Ros. where Ros is the CL contour of the EGRET detection.," considered for a BL Lac to be physically associated with an EGRET source is that the angular distance between the best estimated position of the pair does not exceed $2 R_{95}$, where $R_{95}$ is the CL contour of the EGRET detection."
+ Their claim was based on a positional correlation analysis (using the doubled size for EGRET sources) between the Third EGRET Catalog (3EG. Hartman et al.," Their claim was based on a positional correlation analysis (using the doubled size for EGRET sources) between the Third EGRET Catalog (3EG, Hartman et al."
+ 1999) and the objects identified as BL Lac in the Veron-Cetty Veron (2000) Catalog., 1999) and the objects identified as BL Lac in the Veron-Cetty Veron (2000) Catalog.
+ This results in 14 coincidences. 4 of which are further found to be part of the 5 BL Lacs located within 2.5° of UHECRSs discussed above.," This results in 14 coincidences, 4 of which are further found to be part of the 5 BL Lacs located within $^\circ$ of UHECRs discussed above."
+ The typical Ros radius for EGRET sources is 0.5-1., The typical $R_{95}$ radius for EGRET sources is $^\circ$.
+ Because of such large uncertainties. a standard practice in 7-ray studies aming to give preliminary associations. between EGRET sources and possible counterparts is to study. in addition to the object being proposed. any other coincident system able to generate photons in the EGRET range (100 MeV-10 GeV).," Because of such large uncertainties, a standard practice in $\gamma$ -ray studies aiming to give preliminary associations between EGRET sources and possible counterparts is to study, in addition to the object being proposed, any other coincident system able to generate photons in the EGRET range (100 MeV–10 GeV)."
+ All of the latter should be discarded as the origin of the high energy radiation in order for the association claim to persist., All of the latter should be discarded as the origin of the high energy radiation in order for the association claim to persist.
+ This process usually involves theoretical modelling and multiwavelength observations (see e.g. Caraveo 2002. Reimer et al.," This process usually involves theoretical modelling and multiwavelength observations (see e.g. Caraveo 2002, Reimer et al."
+ 2001. Torres et al.," 2001, Torres et al."
+ 2003a. and references therein).," 2003a, and references therein)."
+"always satisfied, and thus the eccentricity will always be damped.","always satisfied, and thus the eccentricity will always be damped."
+ Thus the observed synchronous binary population can exist in the equilibrium without growth in eccentricity., Thus the observed synchronous binary population can exist in the equilibrium without growth in eccentricity.
+" If the observed synchronous population is assumed to be in this joint opposing evolutionary equilibrium state, then A=1 and Eqn."," If the observed synchronous population is assumed to be in this joint opposing evolutionary equilibrium state, then $A = 1$ and Eqn."
+" [5| is solved for the three unknown quantities (D, Q, and E, ): Table [1] lists and Fig."," \ref{eqn:semimajoraxisevolution} is solved for the three unknown quantities $B$ , $Q$ , and $k_p$ ): Table \ref{tab:BQk} lists and Fig."
+" |l] plots BQ/k, for each of the known synchronous binary systems using observational data (Pravecetal.|/2006;[Pravec&Harris|2007).."," \ref{fig:BQkPlot} plots $B Q / k_p$ for each of the known synchronous binary systems using observational data \citep{Pravec2006, Pravec2007b}."
+" The tidal dissipation number ( is an intensive property that we expect to be similar for all of these bodies, and for small bodies has been estimated to be Q=10? (Goldreich& Soter|1966).."," The tidal dissipation number $Q$ is an intensive property that we expect to be similar for all of these bodies, and for small bodies has been estimated to be $Q = 10^2$ \citep{Goldreich1966}. ."
+ The scatter and size dependence in Fig., The scatter and size dependence in Fig.
+" 1| should be from B and k,, respectively."," \ref{fig:BQkPlot} should be from $B$ and $k_p$, respectively."
+ The BYORP coefficient B does not depend on size; B does depend on the shape of the secondary., The BYORP coefficient $B$ does not depend on size; $B$ does depend on the shape of the secondary.
+ Asteroid shapes can vary greatly introducing scatter in the BYORP coefficient., Asteroid shapes can vary greatly introducing scatter in the BYORP coefficient.
+" estimate B=10? from the shape model of the secondary of 1999 KW,, which does not appear symmetric."," \citet{McMahon2010b} estimate $B = 10^{-3}$ from the shape model of the secondary of 1999 $_4$, which does not appear symmetric."
+" The BYORP coefficient may vary over a few orders of magnitude, especially towards smaller values corresponding to secondaries that are more symmetric."," The BYORP coefficient may vary over a few orders of magnitude, especially towards smaller values corresponding to secondaries that are more symmetric."
+" The tidal Love number may have a dependance on size, and Goldreich&Sari|(2009) predict that the tidal Love number k,,=10? for a “rubble pile"" internal structure."," The tidal Love number may have a dependance on size, and \citet{Goldreich2009} predict that the tidal Love number $k_p = 10^{-5} R_p$ for a “rubble pile” internal structure."
+ In the top plot of Fig. m," In the top plot of Fig. \ref{fig:BQkPlot},"
+" the solid line plots a simple theoretical model of BQ/k, using the estimates of each value from above, so that BQ/k,=10*R71,"," the solid line plots a simple theoretical model of $B Q / k_p$ using the estimates of each value from above, so that $BQ/k_p = 10^4 R_p^{-1}$."
+ Fitting the proportionality constant of the Goldreich&Sari|(2009) tidal Love numbermodel does not significantly changethe results., Fitting the proportionality constant of the \citet{Goldreich2009} tidal Love numbermodel does not significantly changethe results.
+ Dashed lines indicate scatter from the BYORPcoefficient (one order of magnitude larger and two, Dashed lines indicate scatter from the BYORPcoefficient (one order of magnitude larger and two
+11.1 ? with à peak between | and 3 MK.,4.1 $^{-3}$ with a peak between 1 and 3 MK.
+ No indications for a dominant hotcomponent (£> ΜΕ) were found., No indications for a dominant hotcomponent $T \ga$ 4 MK) were found.
+ We present additional evidence for the lack of a solar-type FIP-effect. confirming earlier EUVE results.," We present additional evidence for the lack of a solar-type FIP-effect, confirming earlier EUVE results."
+The courparison of the observed rotation rate of oulsars and stellar models iu the pre-supernova stages indicate that amost stars are losing more angular nomenti fhan currently predicted (οσαetal.2000.IDyschietal. 2001)).,"The comparison of the observed rotation rate of pulsars and stellar models in the pre-supernova stages indicate that most stars are losing more angular momentum than currently predicted \cite{HegerLW00,HMMXII}) )."
+ Normally. the conservation of the centra aneular momentum of a presupernova model would cad to a neutron star spiuuing with a period of 0.1 us. which is about two orders of magnitude faster han he estimate for the most rapid pulsars at birth.," Normally, the conservation of the central angular momentum of a presupernova model would lead to a neutron star spinning with a period of 0.1 ms, which is about two orders of magnitude faster than the estimate for the most rapid pulsars at birth."
+ The question las arisen whether some rotational iustabilities nay play a role iu dissipating the aneulay momenta., The question has arisen whether some rotational instabilities may play a role in dissipating the angular momentum.
+ We cau think iu particular of the Colreich-Sclubert-Fricke (CSF) lustahility (CGoldreich&Schubert1967.Fricke LOGS)). which was a neglieible effect iu the Alain phase and which may play some role in the We and more advanced phases (Ilegeretal. 20001). in particular when there is a very steep O eracicut at the edge of the central deuse core.," We can think in particular of the Golreich-Schubert-Fricke (GSF) instability \cite{GoldreichS67,Fricke68}) ), which has a negligible effect in the Main--Sequence phase and which may play some role in the He--burning and more advanced phases \cite{HegerLW00}) ), in particular when there is a very steep $\Omega$ –gradient at the edge of the central dense core."
+ This instability is eenerallv not accounted for in stellar inodeliug., This instability is generally not accounted for in stellar modeling.
+ The ain of this article is to examine whether the GSE instability is müportant in the pre-superuova stages. when account is given to the effect of the horizontal turbulence in rotating stars which reduces the stabilizing effects of the Hosradieut.," The aim of this article is to examine whether the GSF instability is important in the pre-supernova stages, when account is given to the effect of the horizontal turbulence in rotating stars which reduces the stabilizing effects of the $\mu$ –gradient."
+ Sect., Sect.
+ 2 recalls the basic properties of the CSF instability. Sect.," 2 recalls the basic properties of the GSF instability, Sect."
+ 3 those of the horizontal turbulence., 3 those of the horizontal turbulence.
+ The effects of turbulence ou the CSF instability are examined in Sect., The effects of turbulence on the GSF instability are examined in Sect.
+ L, 4.
+ Sect., Sect.
+ 5 show the results of the ummerical models., 5 show the results of the numerical models.
+ Sect., Sect.
+ 6 eives the conclusion., 6 gives the conclusion.
+ A rotating star with a distribution of the specific angular momentum jf decreasing outwards is subject to the RavleighTavlor iustabilitv: an upward displaced fluid ement will have a higher j than the ambient medium aud is it will continue to move outwards., A rotating star with a distribution of the specific angular momentum $j$ decreasing outwards is subject to the Rayleigh–Taylor instability: an upward displaced fluid element will have a higher $j$ than the ambient medium and thus it will continue to move outwards.
+ In radiative stable uedia. the density stratification has a stabiliziug effect. which may counterbalance the iustabilitv resulting from 16 outwards decrease of j.," In radiative stable media, the density stratification has a stabilizing effect, which may counterbalance the instability resulting from the outwards decrease of $j$."
+ In this respect. the jj exadieu resulting from unclear evolution has a strong stabilizing ffect.," In this respect, the $\mu$ –gradient resulting from nuclear evolution has a strong stabilizing effect."
+ The stability condition is usually expressed by 1e SolbergToiland criterion. giveu in the first part of Eq. (1)).," The stability condition is usually expressed by the Solberg–Hoiland criterion, given in the first part of Eq. \ref{gsf}) )."
+ The CSF instability occurs when the heat diffusion w the fid elements reduces the stabilizing effec of the eutropv stratification in the radiative lavers., The GSF instability occurs when the heat diffusion by the fluid elements reduces the stabilizing effect of the entropy stratification in the radiative layers.
+ The account of a finite viscosity rr together with hermal diffusivity A influences the instability criteria (Fricke1968.Acheson L978)).," The account of a finite viscosity $\nu$ together with thermal diffusivity $K$ influences the instability criteria \cite{Fricke68,Acheson78}) )."
+ These authors found iustabilitv for each of the two conditions where NZ2)7 is. the adiabatic. thermal term of. the Brunt Vaiisalla (BV) frequency and V6 the rotational contribution to DV for au augular velocity O. The viscosity p=(1/23)e6 represents amy source of viscosity. includiug turbulence.," These authors found instability for each of the two conditions where $N^2_{T, \, \mathrm{ad}}$ is the adiabatic thermal term of the Brunt– Väiisällä (BV) frequency and $N^2_{\Omega}$ the rotational contribution to BV for an angular velocity $\Omega$, The viscosity $\nu=(1/3) v \ell$ represents any source of viscosity, including turbulence."
+ a is the distance to the rotation axis and : the vertical coordinate parallel to the rotation axis., $\varpi$ is the distance to the rotation axis and $z$ the vertical coordinate parallel to the rotation axis.
+ The thermal diffusivitv Ads where he various quantities have their usual micauine., The thermal diffusivity $K$ is where the various quantities have their usual meaning.
+"The dust mass of the current BCD sample is estimated by the (90 uum) and (140 um) fluxes, which are suitable to trace the total mass of large grains which have the dominant contribution to the total dust content (Gallianoetal.2005).","The dust mass of the current BCD sample is estimated by the $90~\mu$ m) and $140~\mu$ m) fluxes, which are suitable to trace the total mass of large grains which have the dominant contribution to the total dust content \citep{galliano05}."
+".? The dust mass Ma is related to the flux as citealthildebrand83;; H08) where &, is the mass absorption coefficient of dust grains.", The dust mass $M_\mathrm{d}$ is related to the flux as \\citealt{hildebrand83}; H08) where $\kappa_\nu$ is the mass absorption coefficient of dust grains.
+" The mass absorption coefficient is related with the absorption efficiency Q, as HH08) where s is the grain material density.", The mass absorption coefficient is related with the absorption efficiency $Q_\nu$ as H08) where $s$ is the grain material density.
+" If the grain radius is much smaller than the wavelength, Q,/a is independent of a (Hildebrand1983)."," If the grain radius is much smaller than the wavelength, $Q_\nu /a$ is independent of $a$ \citep{hildebrand83}."
+". The values adopted in Section ?? indicate Q,/a=1.14x10?pm! for silicate and 1.22x10?pm! for graphite at λ=125 jum. By adopting typical densities for silicate and graphite as s=3.3g cm? and 2.25 g cm”, respectively, we obtain the mass absorption coefficients at \=125 wm as K1»5=25.9 and 41.0 cm? g! for silicate and graphite, respectively."," The values adopted in Section \ref{subsec:properties} indicate $Q_\nu /a=1.14\times 10^{-2}~\mu\mathrm{m}^{-1}$ for silicate and $1.23\times 10^{-2}~\mu\mathrm{m}^{-1}$ for graphite at $\lambda =125~\mu$ m. By adopting typical densities for silicate and graphite as $s=3.3$g $^{-3}$ and 2.25 g $^{-3}$, respectively, we obtain the mass absorption coefficients at $\lambda =125~\mu$ m as $\kappa_{125}=25.9$ and 41.0 $^2$ $^{-1}$ for silicate and graphite, respectively."
+ HO8 derived &125—18.8 cm? gt based on Hildebrand(1983)., H08 derived $\kappa_{125}=18.8$ $^2$ $^{-1}$ based on \citet{hildebrand83}.
+". Following H08, we adopt the following expression for &, to estimate the dust mass from the data: The absorption coefficient adopted in Section ?? cannot be fitted with a single power law, but the index lies between 8—1 and 2."," Following H08, we adopt the following expression for $\kappa_\nu$ to estimate the dust mass from the data: The absorption coefficient adopted in Section \ref{subsec:properties} cannot be fitted with a single power law, but the index lies between $\beta =1$ and 2."
+" Thus, if we adopt the wavelength dependence described in Section ??,, the estimated dust mass should be between the dust masses obtained with 6=1 and 6=2."," Thus, if we adopt the wavelength dependence described in Section \ref{subsec:properties}, the estimated dust mass should be between the dust masses obtained with $\beta =1$ and $\beta =2$."
+" We prefer to adopt this simple power-law expression for &, here, since it can be directly compared with the results in HO8 and the simple expression is useful for observational usage."," We prefer to adopt this simple power-law expression for $\kappa_\nu$ here, since it can be directly compared with the results in H08 and the simple expression is useful for observational usage."
+" In Table 2,, we list the dust mass estimated from equation (9))."," In Table \ref{tab:mdust}, we list the dust mass estimated from equation \ref{eq:mdust}) )."
+" We determine the dusttemperature Ty by fitting Av?B,(Ta) (A is a constant) to the 90 jm and 140 jm fluxes.", We determine the dusttemperature $T_{\rm d}$ by fitting $A\nu^{\beta}B_\nu (T_{\rm d})$ $A$ is a constant) to the 90 $\mu$ m and 140 $\mu$ m fluxes.
+ Note that A can be eliminated if we take the flux ratio of those two bands., Note that $A$ can be eliminated if we take the flux ratio of those two bands.
+" The estimated dust temperatures are listed in Table 2,, where T4(8=1) and T4(8=2) denote the dust temperatures evaluated with 3=1 and 2, respectively."," The estimated dust temperatures are listed in Table \ref{tab:mdust}, where $T_\mathrm{d}(\beta =1)$ and $T_\mathrm{d}(\beta =2)$ denote the dust temperatures evaluated with $\beta =1$ and 2, respectively."
+" For the mass absorption coefficient, we adopt Ki25=25.9 cm? gttthe value for silicate), but the dust mass is simply proportional to if another value of &125 is adopted."," For the mass absorption coefficient, we adopt $\kappa_{125}=25.9$ $^2$ $^{-1}$the value for silicate), but the dust mass is simply proportional to $\kappa_{125}^{-1}$ if another value of $\kappa_{125}$ is adopted."
+" The distance is taken from KipsHOS8 for the common sample, and is estimated by using the Galactocentric velocity taken from the NASA/IPAC Extragalactic Database by assuming a Hubble constant of 75 km s! Mpc"," The distance is taken from H08 for the common sample, and is estimated by using the Galactocentric velocity taken from the NASA/IPAC Extragalactic Database by assuming a Hubble constant of 75 km $^{-1}$ $^{-1}$."
+" The dust masses estimated with 3=1 and 2 are denoted as Ma(8!.=1) and Ma(B=2), respectively."," The dust masses estimated with $\beta =1$ and 2 are denoted as $M_\mathrm{d}(\beta =1)$ and $M_\mathrm{d}(\beta =2)$, respectively."
+" Although the complex dependence of quantities in equation (9)) makes it hard to obtain an analytical estimate of the uncertainty in Ma, we have confirmed that the uncertainty is around a factor of 2 by varying the measured fluxes within the errors."," Although the complex dependence of quantities in equation \ref{eq:mdust}) ) makes it hard to obtain an analytical estimate of the uncertainty in $M_\mathrm{d}$ , we have confirmed that the uncertainty is around a factor of 2 by varying the measured fluxes within the errors."
+" From Table 2,, we observe that Ma(3=1) is systematically smaller than Ma(G=2) by factor of ~ 2.5, mainly because of higher dust temperature for a6= 1."," From Table \ref{tab:mdust}, , we observe that $M_\mathrm{d}(\beta =1)$ is systematically smaller than $M_\mathrm{d}(\beta =2)$ by a factor of $\sim 2.5$ , mainly because of higher dust temperature for $\beta =1$ ."
+the polynomial does not fit short-term variations and real term events such as transits.,the polynomial does not fit short-term variations and real short-term events such as transits.
+" We thus write Fluxwhere JD is the Julian date (normalized to range x 1) and a, b, c, and d are the fit parameters for the third degree polynomial."," We thus write where JD is the Julian date (normalized to range $-1 \leq JD \leq 1$ ) and a, b, c, and d are the fit parameters for the third degree polynomial."
+" At the end of this procedure, we have a detrended light curve per filter for each star."," At the end of this procedure, we have a detrended light curve per filter for each star."
+" After this step, proceeds in removing the jumps."," After this step, proceeds in removing the jumps."
+" To identify the cleanest light curve for a reliable jump removal, we create ""sub-light curves"", which have a typical duration of a day."," To identify the cleanest light curve for a reliable jump removal, we create ”sub-light curves”, which have a typical duration of a day."
+" Thus, for the IRa01 field we create 60 ""sublight"" curves, called simply light curves in the following."," Thus, for the IRa01 field we create 60 ”sublight” curves, called simply light curves in the following."
+ These 60 blocks were selected after we checked various combinations., These 60 blocks were selected after we checked various combinations.
+" If the number of blocks is too small, the probability of including a jump in the ""sublight"" curve increases."," If the number of blocks is too small, the probability of including a jump in the ”sublight” curve increases."
+ Figure 3 shows the optimal block number versus standard deviation., Figure \ref{bloc} shows the optimal block number versus standard deviation.
+" We assume that there are three full light curves for a given star in each band with N points per light curve which we denote by Fri, Fgi, and Fg; with i=1,N the individual data values in the red, green, and blue filters, respectively."," We assume that there are three full light curves for a given star in each band with $N$ points per light curve which we denote by $F_{R,i}$, $F_{G,i}$, and $F_{B,i}$ with $i=1,N$ the individual data values in the red, green, and blue filters, respectively."
+ We then divide each color light curve into 60 sub-light curves (one sub-light curve per day for IRa01 - 60 days)., We then divide each color light curve into 60 sub-light curves (one sub-light curve per day for IRa01 - 60 days).
+" For each sub-light curve, we calculate the mean value MR, MG, and MB and normalize each sub-light curve by its mean value."," For each sub-light curve, we calculate the mean value $MR$, $MG$, and $MB$ and normalize each sub-light curve by its mean value."
+" We then compute new, normalized sub-light curves NF to be for each filter band, and it is clear that all of these light curves have a mean of unity."," We then compute new, normalized sub-light curves $NF$ to be for each filter band, and it is clear that all of these light curves have a mean of unity."
+" This normalization is necessary otherwise the entire process would be dominated by the light curve of the strongest signal, which is usually the red light curve."," This normalization is necessary otherwise the entire process would be dominated by the light curve of the strongest signal, which is usually the red light curve."
+" As a side effect, normalizes the depth of a possible transit in all filters using Eq. 2,,"," As a side effect, normalizes the depth of a possible transit in all filters using Eq. \ref{eq2},"
+" so when the algorithm continues with its next steps, all transit events in each filter will have the same depth,thus does not destroy real signals from the transits."," so when the algorithm continues with its next steps, all transit events in each filter will have the same depth,thus does not destroy real signals from the transits."
+" The normalized light curves now have the same mean, but their dispersions, differ."," The normalized light curves now have the same mean, but their dispersions, differ."
+" Our next goal is to identify the instrumental scatter caused, for example by jumps in each light curve and differentiate this instrumental scatter from statistical noise."," Our next goal is to identify the instrumental scatter caused, for example by jumps in each light curve and differentiate this instrumental scatter from statistical noise."
+" To achieve this, extracts five random packages of twenty adjacent points each from all colour bands and calculates the standard deviation of each package per filter; the result should represent a good estimate of the correct light curve value at that time."," To achieve this, extracts five random packages of twenty adjacent points each from all colour bands and calculates the standard deviation of each package per filter; the result should represent a good estimate of the correct light curve value at that time."
+" If we use many packages, the probability of including jumps increases."," If we use many packages, the probability of including jumps increases."
+" The correct combination of packages and points is a function of the duration of the jumps, which is a random value, thus there is no"," The correct combination of packages and points is a function of the duration of the jumps, which is a random value, thus there is no"
+A straight-forward but leugthy derivation shows that equation can be recast ito an equation for c. (Ej Ey pon dass where the coefficient. d; is reht]] | Coa?,"A straight-forward but lengthy derivation shows that equation can be recast into an equation for $\psi_0$, _1 - _0 + (d_1 - d_2) _0 = 0, where the coefficient $d_i$ is d_i = ] +. . - ."
+ Uere |J! refers to the power-law index for pans. or that for p near the surface.," Here $\beta$ refers to the power-law index for $\rho_{\rm surf}$ , or that for $\rho$ near the surface."
+ The benefit of the transformation introduced in equation is that d; contaius no derivatives on c., The benefit of the transformation introduced in equation is that $d_i$ contains no derivatives on $\psi$.
+ Near the surface. (X—const. d;=(0 as is expected and equation is separable.," Near the surface, $X = {\rm const}$, $d_i = 0$ as is expected and equation is separable."
+ But in deeper region of the planet. .X is a complex fiction of .rj and c» and equation is not separable.," But in deeper region of the planet, $X$ is a complex function of $x_1$ and $x_2$ and equation is not separable."
+ However. if Wis a smoothlv varviug function with a scale length being the radius of theplanct!.. one can show that. in the WISB region. djcy~cy and is 1/A2 sinaller than the £;coy term.," However, if $X$ is a smoothly varying function with a scale length being the radius of the, one can show that, in the WKB region, $d_i\psi_0 \sim
+\psi_0$ and is $\sim 1/\lambda^2$ smaller than the ${\cal E}_i \psi_0$ term."
+ So if we ignore the d; terms. we oulv introduce errors of order OA3}κ1.," So if we ignore the $d_i$ terms, we only introduce errors of order ${\cal O}(\lambda^{-2}) \ll 1$."
+ It is worth pointing out that the uou-WB region occurs ucar the surface where our solution is exact., It is worth pointing out that the non-WKB region occurs near the surface where our solution is exact.
+" Iu conclusion. we cau adopt the following approximate solution for c U—VXv with ος) satisfiug ὅτε,|We;=0."," In conclusion, we can adopt the following approximate solution for $\psi$ = _0 = _1(x_1) _2(x_2), with $\psi_i(x_i)$ satisfying ${\cal E}_i \psi_i + K^2 \psi_i = 0$."
+ This solution is exact near the surface aud is accurate to Q(1/A3) in the WISB region.," This solution is exact near the surface and is accurate to ${\cal
+O}(1/\lambda^2)$ in the WKB region."
+ The latter attribute indicates that the approximate solution describes accurately both the euvelope and the phase of the actual inertialanode cieenfuuction., The latter attribute indicates that the approximate solution describes accurately both the envelope and the phase of the actual inertial-mode eigenfunction.
+ Having discussed odinethods to obtain inertialanocdoe eieenfuuctious dn. various density profiles. here we focus on studyiug general properties of inertial-niodes.," Having discussed methods to obtain inertial-mode eigenfunctions in various density profiles, here we focus on studying general properties of inertial-modes."
+ This include the WEKD properties the nonualization relatiouship aud density of modes in the frequency spectrum.," This include the WKB properties, the normalization relationship and density of modes in the frequency spectrum."
+ Moreover. we attempt to eive readers a eraplical impression of how inertialimodes look like both inside aud outside the planet.," Moreover, we attempt to give readers a graphical impression of how inertial-modes look like both inside and outside the planet."
+ Lasthv. we compare inertialinodes against well known eravitv- aud pressure-modes to eain intuition for this brauch of eigeumodes.," Lastly, we compare inertial-modes against well known gravity- and pressure-modes to gain intuition for this branch of eigenmodes."
+ We first derive the WIXB dispersion relation for inertialinodes., We first derive the WKB dispersion relation for inertial-modes.
+ We then determine the confue of the WISB region. and end with a eeueral derivation for the WISB cuvelope of iuertial-iiode amplitude.," We then determine the confine of the WKB region, and end with a general derivation for the WKB envelope of inertial-mode amplitude."
+ Tn the WISB region. let. V.z;k aud 0/0:zik. equation yields. DO qud or pzmkefhe <1.," In the WKB region, let $\bnabla \approx i \boldk $ and $\partial/\partial z
+\approx i k_z$, equation yields, k^2 - q^2 k_z^2 0, or $\mu \approx {{k_z}/{k}} \leq 1$ ."
+ This is to be compared with result from a more careful derivation 2.2.0)) which shows przsf(nemA).," This is to be compared with result from a more careful derivation \ref{subsubsec:dispersion}) ) which shows $\mu
+\approx sin(n_2 \pi/\lambda)$."
+ Since the ο axis is largely. along the + axis (Fig. 8)). ke~ofHR.," Since the $x_2$ axis is largely along the $z$ axis (Fig. \ref{fig:x1x2}) ), $k_z \sim
+n_2/R$."
+ So we have bh~ASR—2(n|no)R. with R heineg the planet radius (normalized to be one).," So we have $k \sim \lambda/R \sim 2 (n_1 + n_2)/R$, with $R$ being the planet radius (normalized to be one)."
+ Such a dispersion relation iuples that the mode frequeney (w=204) does not depend ou the ummber of wigeles iu a mode. but rather ou the direction of wavepropagation.," Such a dispersion relation implies that the mode frequency $\omega = 2\Omega \mu$ ) does not depend on the number of wiggles in a mode, but rather on the direction of wavepropagation."
+ Modes that propagate close to the rotation axis have higher frequencies 20) than those that propagate close to the equator (uw~ 00.," Modes that propagate close to the rotation axis have higher frequencies $\omega
+\sim 2 \Omega$ ) than those that propagate close to the equator $\omega \sim 0$ )."
+ Such a dispersion relation can be understood by the following physical argument., Such a dispersion relation can be understood by the following physical argument.
+" We relate € to c using equation(13).. So the fluid velocity (ο= 0£/0f) is perpendicular to the phase velocity (ey=οRP R). ov wv, oO."," We relate $\boldxi$ to $\psi$ using equation, i. So the fluid velocity $\boldv = \partial \boldxi/\partial t$ ) is perpendicular to the phase velocity $\boldv_{\rm ph} = \omega/k^2
+\boldk$ ), or $\boldv
+\cdot \boldv_{\rm ph} \approx 0$ ."
+ Tuertial-waves are larecly transverse waves., Inertial-waves are largely transverse waves.
+ A mode with its phase velocity aloug the equator will show fluid motion in the : direction., A mode with its phase velocity along the equator will show fluid motion in the $z$ direction.
+ As a result. it experieuces little Coriolis force. and has a frequency that is close to zero.," As a result, it experiences little Coriolis force, and has a frequency that is close to zero."
+ Iu. coutrast.a modewith & along the : direction will experience the strongest restoriug force and will have the highest frequelcy.," In contrast,a modewith $\boldk $ along the $z$ direction will experience the strongest restoring force and will have the highest frequency."
+ While the phase velocity of theinertia-awave vpn wfkk. the group velocity⋅ of. the inertial-xvave. runsas | = en," While the phase velocity of theinertial-wave ${\boldv}_{\rm ph} =
+\omega/k^2 \boldk$ , the group velocity of the inertial-wave runsas = = + = ( -"
+"UUsing equations (38), (39) & (52) we get, wwhere q is the deceleration parameterdefined as 2—22.","Using equations (38), (39) $\&$ (52) we get, where $q $ is the deceleration parameterdefined as $ - \frac{\ddot{R}R}{\dot{R}^{2}}$."
+ For accelerating model the value of q should be negative., For accelerating model the value of $q$ should be negative.
+" So for real value Us, xa+$(m?d(d--1)2(1—dm))>(9—dm) which gives xagTaama»ηd(1)4-2(1—dm))>|q |."," So for real value of $v_{s}$, $\frac{1 + \alpha}{d + 3} + \frac{\alpha}{k}\left\{m^{2}d(d + 1) +
+2(1 - dm)\right\} > \frac{2\alpha \mid q \mid}{k}(2 - dm)$ which gives $\frac{k(1 + \alpha)}{2 \alpha (d + 3)(2 - dm)} +
+\frac{1}{2(2-dm)}\left\{m^{2}d(d + 1) + 2(1-dm)\right\} > \mid q
+\mid$ ."
+" This relation gives an upperbound for the g, which may Unihave observational consequences."," This relation gives an upperbound for the $q$, which may have observational consequences."
+" Since 0-«a«1, it can not exceed the speed of light."," Since $0<\alpha<1$, it can not exceed the speed of light."
+ But for larger values of a the speed of sound exceeds the speed of light., But for larger values of $\alpha$ the speed of sound exceeds the speed of light.
+ It also depends upon the value of m., It also depends upon the value of $m$.
+ In cosmology a speed of sound exceeding the speed of light does not contradict the causality [0].., In cosmology a speed of sound exceeding the speed of light does not contradict the causality \cite{Gorini}.
+" LLet a=0.5, m= ᾱ, and d=6, therefore k should be 2 and M equal to 38 in this case."," Let $\alpha = 0.5$, $m = \frac{1}{6}$ , and $d = 6$, therefore $k$ should be $\frac{5}{6}$ and $M$ equal to $ \frac{88}{27}$ in this case."
+" If we calculate v?, we get two conditions (i) when |q|<0.923 then v,«c, (ii) otherwise Us> C."," If we calculate $v_{s}^{2}$, we get two conditions (i) when $\mid q \mid < 0.923 $ then $v_{s} < c$, (ii) otherwise $v_{s} > c$ ."
+" FFollowing the analysis givenin the previous section we get the analogous energy density poand pressure p, for the scalar field as", Following the analysis givenin the previous section we get the analogous energy density $\rho_{\phi}$and pressure $p_{\phi}$ for the scalar field as
+The first burst of star formation in the Universe is thought to give rise to massive stars. the so-called Population HL with current theory predicting masses in the range 20-150 M. (Abeletal..2002:Bromm2002;O'Shea&Norman.2007;Yoshidaetal.. 2008)..,"The first burst of star formation in the Universe is thought to give rise to massive stars, the so-called Population III, with current theory predicting masses in the range 20-150 $_{\odot}$ \citep{2002Sci...295...93A, 2002ApJ...564...23B, 2007ApJ...654...66O, 2008Sci...321..669Y}."
+ This contrasts with present-day star formation. which tends to yield stars with masses less than | Μι (Kroupa.2002:Chabrier.2003).. and so at some point in the evolution of the Universe there must have been a transition from primordial (Pop.," This contrasts with present-day star formation, which tends to yield stars with masses less than 1 $_{\odot}$ \citep{2002Sci...295...82K, 2003PASP..115..763C}, and so at some point in the evolution of the Universe there must have been a transition from primordial (Pop."
+ HL) star formation to the mode of star formation we see today (Pop., III) star formation to the mode of star formation we see today (Pop.
+ When gas collapses to form stars. gravitational energy is transformed to thermal energy and unless this can be dissipated in some fashion. it will eventually halt the collapse.," When gas collapses to form stars, gravitational energy is transformed to thermal energy and unless this can be dissipated in some fashion, it will eventually halt the collapse."
+ Thermal energy can be dissipated by processes such as atomic fine structure line emission. molecular rotational or vibrational line emission. or the heating of dust grains.," Thermal energy can be dissipated by processes such as atomic fine structure line emission, molecular rotational or vibrational line emission, or the heating of dust grains."
+ In some cases. these processes are able to cool the gas significantly during the collapse.," In some cases, these processes are able to cool the gas significantly during the collapse."
+ This temperature drop can promote gravitational fragmentation (MaeLow&Klessen.2004.Bonnelletal..2007) by diminishing the Jeans mass. which means that instead of forming very massive clumps. with fragment masses corresponding to the initial Jeans mass in the cloud. it can instead form even more fragments with lower If the gas is cooled only by molecular hydrogen emission. numerical simulations show that the stars should be very massive (Abeletal..2002:BrommO'Shea&Norman.2007;Yoshidaetal.. 2008)..," This temperature drop can promote gravitational fragmentation \citep{2004RvMP...76..125M, 2007prpl.conf..149B} by diminishing the Jeans mass, which means that instead of forming very massive clumps, with fragment masses corresponding to the initial Jeans mass in the cloud, it can instead form even more fragments with lower If the gas is cooled only by molecular hydrogen emission, numerical simulations show that the stars should be very massive \citep{2002Sci...295...93A, 2002ApJ...564...23B, 2007ApJ...654...66O, 2008Sci...321..669Y}."
+ This happens because the H» cooling becomes inefficient for temperatures bellow 200K and densities above 10'em7*., This happens because the $_2$ cooling becomes inefficient for temperatures bellow 200K and densities above $10^4 \rm cm^{-3}$.
+" At this temperature and density. the mean Jeans mass at cloud fragmentation ts 1.000 times larger than in present-day molecular clouds. for an atomic gas with temperature 754, and number density Πω. Metal line cooling and dust cooling are effective at lower temperatures and larger densities. and so the most widely accepted cause for the transition from Pop."," At this temperature and density, the mean Jeans mass at cloud fragmentation is 1,000 times larger than in present-day molecular clouds, for an atomic gas with temperature $T_{\rm frag}$ and number density $n_{\rm frag}$ Metal line cooling and dust cooling are effective at lower temperatures and larger densities, and so the most widely accepted cause for the transition from Pop."
+ Π to Pop., III to Pop.
+ Η is metal enrichment of the interstellar medium by the previous generations of stars., II is metal enrichment of the interstellar medium by the previous generations of stars.
+ This suggests that there might be a critical metallicity Zi at which the mode of star formation The main coolants that have been studied in the literature are CII and OI fine structure emission (Brommetal..2001:etal. 2007).. and dust emission.," This suggests that there might be a critical metallicity $_{\rm crit}$ at which the mode of star formation The main coolants that have been studied in the literature are CII and OI fine structure emission \citep{2001MNRAS.328..969B, 2003Natur.425..812B, 2006ApJ...643...26S, 2007MNRAS.380L..40F}, and dust emission."
+ C and O are identified as the key species because in the temperature and density conditions that characterise the early phases of Pop., C and O are identified as the key species because in the temperature and density conditions that characterise the early phases of Pop.
+ IL star formation. the OI and CII fine-structure lines dominate over all other metal transitions (Hollenbach&McKee.1989)..," III star formation, the OI and CII fine-structure lines dominate over all other metal transitions \citep{1989ApJ...342..306H}."
+ By equating the CII or OI fine structure cooling rate to the compressional heating rate due to free-fall collapse. one can define critical abundances [Ο/Η]=—3.5 and [O/H]=-3.0! for efficient metal line cooling (Bromm&Loeb.2003)..," By equating the CII or OI fine structure cooling rate to the compressional heating rate due to free-fall collapse, one can define critical abundances $[\rm C/H] = -3.5$ and $[ \rm O/H] = -3.0$ for efficient metal line cooling \citep{2003Natur.425..812B}."
+ However. previous works (Jappsenetal.2009a.b) show that this metallicity threshold does not represent a critical metallicity: the fact that metal-line cooling has a larger value than the compressional heating does not necessarily lead to," However, previous works \citep{2009ApJ...696.1065J, 2009ApJ...694.1161J} show that this metallicity threshold does not represent a critical metallicity: the fact that metal-line cooling has a larger value than the compressional heating does not necessarily lead to"
+branch and top of the phase-fFolded. X-ray. lighteurve etal. 2003).,branch and top of the phase-folded X-ray lightcurve \citep{vilhu1}.
+ This pointing occurred 2.2 cays after the peak of a major lave (713 Jv in the 15 CGllz band) when the source was in the FSAR state., This pointing occurred 2.2 days after the peak of a major flare $\sim$ 13 Jy in the 15 GHz band) when the source was in the FSXR state.
+ We note the presence of one 8.5 mllz (~ 120 $3) QPO in this pointing. occurring in two separate episodes cach lasting about 12 eveles (data between these wo separate episodes are devoid of QPOs).," We note the presence of one 8.5 mHz $\sim$ 120 s) QPO in this pointing, occurring in two separate episodes each lasting about 12 cycles (data between these two separate episodes are devoid of QPOs)."
+ In the PDS his QPO corresponds to rms and most of its power is concentrated in the three eveles visible at the beginning of the lighteurve in Fig. 2extitb.., In the PDS this QPO corresponds to rms and most of its power is concentrated in the three cycles visible at the beginning of the lightcurve in Fig. \ref{fig-2}.
+ A blow-up of this region is plotted in Fig., A blow-up of this region is plotted in Fig.
+ 3 where there is possible indication of a weaker 30 mllz (~ 30 s) QPO-like period. (seen in the PDS as à multiple-»eaked. region but not. producing a peak above the significance level) ancl LOO mllz ( 10 sec) QDPO-like period (not seen in the PDS. however possibly contributing power only in this short segment of lighteurve) in. alternating rhythm: three times the 30 sec period whieh form the peak of the 120 sec QPO and four times the 10 sec period which form the valley of the 120 sec QPO.," \ref{fig-3} where there is possible indication of a weaker 30 mHz $\sim$ 30 s) QPO-like period (seen in the PDS as a multiple-peaked region but not producing a peak above the significance level) and 100 mHz $\sim$ 10 sec) QPO-like period (not seen in the PDS, however possibly contributing power only in this short segment of lightcurve) in alternating rhythm: three times the 30 sec period which form the peak of the 120 sec QPO and four times the 10 sec period which form the valley of the 120 sec QPO."
+ This pointing occurred 17 days after a major Hare event. (~ 1.5 Jv in the 15 Gllz band)., This pointing occurred 17 days after a major flare event $\sim$ 1.5 Jy in the 15 GHz band).
+ During this pointing the 15 Gllz racio llux censity is ~ 50 my and thesource is in the FLAR state., During this pointing the 15 GHz radio flux density is $\sim$ 50 mJy and thesource is in the FHXR state.
+ Phere is a strong 21 mlz (~ 50s) OPO lasting about 20 eveles corresponding to rms in the PDS., There is a strong 21 mHz $\sim$ 50 s) QPO lasting about 20 cycles corresponding to rms in the PDS.
+ In the PDS there are also two peaks with lower significances that correspond to frequencies 9 and 31 mill., In the PDS there are also two peaks with lower significances that correspond to frequencies 9 and 31 mHz.
+ An in-depth analysis is bevond. the scope of this paper., An in-depth analysis is beyond the scope of this paper.
+ Llere we speculate briellv on the origin of the QPOs found in this study., Here we speculate briefly on the origin of the QPOs found in this study.
+ e Lowe assume that the QPOs are N-rav. emitting blobs or regions in orbit around. the compact. object. the possible timescales are determined. by the innermost stable circular orbit (SCO. highest. frequency) anc the Roche. Lobe radius or the bow shock that is formed. as the compact object plows thestellar wind of the companion (lowest Lrequency)," $\bullet$ If we assume that the QPOs are X-ray emitting blobs or regions in orbit around the compact object, the possible timescales are determined by the innermost stable circular orbit (ISCO, highest frequency) and the Roche Lobe radius or the bow shock that is formed as the compact object plows thestellar wind of the companion (lowest frequency)."
+ Identifving the Weplerian radius of the QPOs (~0.16R..) with the ISCO results in a compact object of thousands of solar masses therefore rendering this scenario rather unlikely., Identifying the Keplerian radius of the QPOs $\sim$ $_{\odot}$ ) with the ISCO results in a compact object of thousands of solar masses therefore rendering this scenario rather unlikely.
+ At the other end. of the length. scale. we finc the Roche Lobe radius to be approximately 0.7.2.1 1k. and the characteristic length. scale. of the bow shock. Res=Εμμπρι~ BA5.7 Ho. with the compac object moving with velocity Veo through the stellar wine of density p.," At the other end of the length scale, we find the Roche Lobe radius to be approximately 0.7–2.1 $_{\odot}$ and the characteristic length scale of the bow shock, $R_{BS} = \dot{m} v_{wind} / 4 \pi \rho V^{2}_{C}\sim$ 3.1–5.7 $_{\odot}$, with the compact object moving with velocity $_{C}$ through the stellar wind of density $\rho$."
+ The above values have been caleulatecl using the mass ratio of the svstem (Alye/Ale = p elven in Vilhuetal. 2009)). orbital period of 4.8 hours. compac object mass Me: = 1.430 AL. and typical parameters for wind accretion from WR. stars (μα = 1600 km J| with velocity law following Springmann1994. and mass-loss rate m= l6. 10M. Following Szostck&Zdziarski2008)).," The above values have been calculated using the mass ratio of the system $_{WR}$ $_{C}$ = $^{+1.7}_{-1.4}$ given in \citealt{vilhu2}) ), orbital period of 4.8 hours, compact object mass $_{C}$ = 1.4–30 $_{\odot}$ and typical parameters for wind accretion from WR stars $v_{wind}$ = 1600 km $^{-1}$ with velocity law following \citealt{springmann} and mass-loss rate $\dot{m}$ = 1.6 $\times$ $^{-5} M_{\odot}$ $^{-1}$ following \citealt{szostek}) )."
+ Therefore. for a reasonable range of compact object masses the QPOs seem not to be associated with these radii," Therefore, for a reasonable range of compact object masses the QPOs seem not to be associated with these radii."
+" eoscillalions? In the literature there are several methods as to how low-lrequency QPOs can be invoked in the aceretion disk. e.g. low-frequeney. ""ονnamo eveles” in the azimuthal magnetic field or Accretion-Ljection Instability (Varnicre&Tageer2002) to name a few."," $\bullet$ In the literature there are several methods as to how low-frequency QPOs can be invoked in the accretion disk, e.g. low-frequency “dynamo cycles"" in the azimuthal magnetic field \citep{oneill} or Accretion-Ejection Instability \citep{varniere} to name a few."
+ Llowever. the geometry of the system night present a challenge to clisk-basecl phenomena.," However, the geometry of the system might present a challenge to disk-based phenomena."
+ Lf the accretion disk is perpendicular to the jet and if the jet is inclined  14 degrees to the line of sight (Mioduszewskietal. 2001).. then we should be viewing the disk face-on.," If the accretion disk is perpendicular to the jet and if the jet is inclined $\sim$ 14 degrees to the line of sight \citep{mio}, then we should be viewing the disk face-on."
+ The oblem then arises as to how the actual modulation can » seen at all., The problem then arises as to how the actual modulation can be seen at all.
+ A further interesting question is how would he simultaneous frequencies observed. in the 2000 Apr 3 »ointing arise in the disk. especially when they seem to be oducing a clear pattern?," A further interesting question is how would the simultaneous frequencies observed in the 2000 Apr 3 pointing arise in the disk, especially when they seem to be producing a clear pattern?"
+ Are they causally linked and/or xocduced by the same underlying mechanism?, Are they causally linked and/or produced by the same underlying mechanism?
+ Interestingly. he 2009 Aue 9 pointing indicates that the peaks in the ο form a 3:2:1 frequency relation. the same relation as jas been reported from Ser A* (Aschenbachetal.2004).," Interestingly, the 2009 Aug 9 pointing indicates that the peaks in the PDS form a 3:2:1 frequency relation, the same relation as has been reported from Sgr A* \citep{aschenbach}."
+. eoscillations? As the QPOs appear o be present in energy. bands from 215 keV and as they appear to arise in flaring radio/X-ray states that can be modelled solely with a Comptonization component. (Ix10) his could imply that the actual component causing the Comptonization of soft photons is oscillating., $\bullet$ As the QPOs appear to be present in energy bands from 2–15 keV and as they appear to arise in flaring radio/X-ray states that can be modelled solely with a Comptonization component (K10) this could imply that the actual component causing the Comptonization of soft photons is oscillating.
+ An oscillating corona due to a magneto-acoustic wave propagating within he corona producing multiple QPOs has been theoretically discussed in Cabanacetal., An oscillating corona due to a magneto-acoustic wave propagating within the corona producing multiple QPOs has been theoretically discussed in \citet{cabanac}. .
+ (2010).. e jet? As the QPOs appear to occur after major radio Hares. the latter most likely signaling a jet ejection event. the two are probably linked.," $\bullet$ As the QPOs appear to occur after major radio flares, the latter most likely signaling a jet ejection event, the two are probably linked."
+ This connection could be two-fold: either the jet ds shadowing. therefore modulating. an underlying oscillation," This connection could be two-fold: either the jet is shadowing, therefore modulating, an underlying oscillation"
+standard deviation is 0.5.,standard deviation is 0.5.
+ This is an miportaut point. for the sample of galaxies that have both cluission lines and MIPS detections. the full star formation rates of the optical galaxies cau be recovered after correcting for the extinction usimg optical enuission lines.," This is an important point, for the sample of galaxies that have both emission lines and MIPS detections, the full star formation rates of the optical galaxies can be recovered after correcting for the extinction using optical emission lines."
+ We sugecstOO that this is because the same dust is causing extinction of the optical lines. as is producing the Mid-IB. eniission.," We suggest that this is because the same dust is causing extinction of the optical lines, as is producing the Mid-IR emission."
+ Afauy studies lave used the {μι enüssion alone as a proxy for total infrared huuinositv (27?3...," Many studies have used the $\mu$ m emission alone as a proxy for total infrared luminosity \citep{rie09,bai06,mah10}."
+ Certainly the SED fitting technique is more robust. as it is based on up to 15 photometric points.," Certainly the SED fitting technique is more robust, as it is based on up to 15 photometric points."
+ That the method is more accurate. is still open for debate. however. receut results from theHerschel space telescope support this.," That the method is more accurate, is still open for debate, however, recent results from the space telescope support this."
+ These new data show how star formation rates based on [jnu data underestimate rates derived by imeasurnus onger waveleneth fluxes. in of the cases (?)..," These new data show how star formation rates based on $\mu$ m data underestimate rates derived by measuring longer wavelength fluxes, in of the cases \citep{raw10}."
+ Our ata also show a region where the lan data underestimates the rate calculated frou the ull SED fitting., Our data also show a region where the $\mu$ m data underestimates the rate calculated from the full SED fitting.
+ In Figure 12.. our results agree with the relatiouship of ? (the dashed liue). out o the high dux end.," In Figure \ref{lirf24}, our results agree with the relationship of \citet{rie09} (the dashed line), out to the high flux end."
+ Bevoud [12s. he tan uuderestimates the TR huuiuositv as nieasured by he hest-fit SED.," Beyond $\,$ mJy, the $\mu$ m underestimates the IR luminosity as measured by the best-fit SED."
+ Iu this section. we show that our results confiii aud extend those of ?.. who have receutly explored the specific star formation inComa cluster galaxies.," In this section, we show that our results confirm and extend those of \citet{mah10}, who have recently explored the specific star formation inComa cluster galaxies."
+ These authors have used the Yan - z/) color as a proxy for quautifviug sSFR., These authors have used the $\mu$ m - $^{\prime}$ ) color as a proxy for quantifying sSFR.
+ Our iethod relies on selecting galaxies based. ou the physical criteria of stellar mass iud SFR., Our method relies on selecting galaxies based on the physical criteria of stellar mass and SFR.
+ The first panel of Figure 13 shows that the usage of [jiu - z/) recovers normal star forming and starbursts. ealaxies with specific SER ο 0.01 iL although. their color cut of -6 also includes contanunation by many lower sSERBR galaxies.," The first panel of Figure \ref{m24z} shows that the usage of $\mu$ m - $^{\prime}$ ) recovers normal star forming and starbursts, galaxies with specific SFR $>$ 0.01 $^{-1}$, although, their color cut of -6 also includes contamination by many lower sSFR galaxies."
+ There is also a population of galaxies with low and normal specific star formation rates (between 0.0005-0.1 Cyr. 1) which lie above the lynn - 21) color cut of -6., There is also a population of galaxies with low and normal specific star formation rates (between 0.0005-0.1 $^{-1}$ ) which lie above the $\mu$ m - $^{\prime}$ ) color cut of -6.
+ Oue of the results of our work is that the infrared star formune and starburst salaxies are blue. having (e - 47) « 0.7.," One of the results of our work is that the infrared star forming and starburst galaxies are blue, having $^{\prime}$ - $^{\prime}$ ) $<$ 0.7."
+ The ceutral pancl of Figure 13 confinis that the (21jiu - z/) > -6 selects the same blue aud red populations based ou the optical color cut., The central panel of Figure \ref{m24z} confirms that the $\mu$ m - $^{\prime}$ ) $>$ -6 selects the same blue and red populations based on the optical color cut.
+ The righ paucl o| Figure 13 shows that we are able to probe deep iuto the low-mass regine. casting ealaxies with stellar masses as low as LOOAL..," The right panel of Figure \ref{m24z} shows that we are able to probe deep into the low-mass regime, measuring galaxies with stellar masses as low as $^{7.5}$ $_{\odot}$."
+ The vertical dotted line shows the limit of the ? oper at a muss of LOSPAL..., The vertical dotted line shows the limit of the \citet{mah10} paper at a mass of $^{8.5}$ $_{\odot}$.
+ As we saw in Fieure 9.. the low mass dwart galaxies are Huportant as the lighest specific sSSER ealaxics. the stavbuasts. all have masses lower than 10? ...," As we saw in Figure \ref{sfrmass}, the low mass dwarf galaxies are important as the highest specific sSFR galaxies, the starbursts, all have masses lower than $^{9}$ $_{\odot}$."
+ Most of the nuage extends bevoud the virial radius of Coma. parametrized by roug. the characteristic radius for the cluster.," Most of the imaging extends beyond the virial radius of Coma, parametrized by $_{200}$, the characteristic radius for the cluster."
+ We can therefore compare the galaxies iu the cluster core. «rogo. o those on the edge of the cluster ceuter. bevonud half of τουυ.," We can therefore compare the galaxies in the cluster core, $\times$ $_{200}$, to those on the edge of the cluster center, beyond half of $_{200}$."
+ For Coma. roo) has been measured to be IMpec (?)..," For Coma, $_{200}$ has been measured to be $^{-1}$ Mpc \citep{gel99}."
+ At the cluster redshift. this trauslates to royy= Ne’.," At the cluster redshift, this translates to $_{200}$ = $\,$$^{\prime}$."
+ Figure |l1 shows the positions of the spectroscopically confirmed [gin emitting cluster members on the sky., Figure \ref{radec} shows the positions of the spectroscopically confirmed $\mu$ m emitting cluster members on the sky.
+ Most davarft galaxies are starbursts or normal star forming galaxies. and few exist iu the cluster core.," Most dwarf galaxies are starbursts or normal star forming galaxies, and few exist in the cluster core."
+ The suppression of star formation in the cluster core is also is consistent with ? who found a dearth of active galaxies in the core of the Coma cluster based ou SDSS spectroscopy., The suppression of star formation in the cluster core is also is consistent with \citet{cas01} who found a dearth of active galaxies in the core of the Coma cluster based on SDSS spectroscopy.
+ We compute the ealaxy deusitv using the positions of 1557 red-sequence ealaxies from. the SDSS with 12/20. musing a 2D wavelet transform as in ?..," We compute the galaxy density using the positions of 1557 red-sequence galaxies from the SDSS with $<$ $^{\prime}$$<$ 20, using a 2D wavelet transform as in \citet{fad98}."
+ The red-sequeuce is the same ax shown in Figure 10.. where (οἱ) = «1 1.320. aud we inchide a range of £0.07 iu (οὐ-o 17) color.," The red-sequence is the same as shown in Figure \ref{grcolmag}, where $^{\prime}$ $^{\prime}$ ) = $\times$ $^{\prime}$ - 1.320, and we include a range of $\pm$ 0.07 in $^{\prime}$ $^{\prime}$ ) color."
+ The units of galaxy deusitv are thus in uuuber of red-sequence galaxies js) per square degree. represeuted as contours in Fieure 11..," The units of galaxy density are thus in number of red-sequence galaxies $_{RS}$ ) per square degree, represented as contours in Figure \ref{radec}. ."
+ The core and svell-studied eroup NGC 1839 at (191.35. 27.5) are easily detected.," The core and well-studied group NGC 4839 at (194.35, 27.5) are easily detected."
+ A previously uudetected eroup at 27.7dog) is also detected along the direction of the filament toward Abell 1367.," A previously undetected group at $\,$ 27.7deg) is also detected along the direction of the filament toward Abell 1367."
+ The top panel of Figure 15 reveals a dearth of starburst galaxies iu the cluster core. where the ealaxv density ngs200.," The top panel of Figure \ref{dens} reveals a dearth of starburst galaxies in the cluster core, where the galaxy density $_{RS}$$ < $ 300."
+ In this fleure. we aim," In this figure, we aim"
+spanned is very wide. covering almost 4+ dex.,"spanned is very wide, covering almost 4 dex."
+" Even if at first we will consider the general field altogether. it is useful to inspect the local density distributions of group. binary and single galaxies separately. as shown in refpm2ge,isto."," Even if at first we will consider the general field altogether, it is useful to inspect the local density distributions of group, binary and single galaxies separately, as shown in \\ref{pm2gc_histo}."
+""" Single galaxies are preferentially located in the lowest density bins. groups in the highest and binaries in the intermediate range."," Single galaxies are preferentially located in the lowest density bins, groups in the highest and binaries in the intermediate range."
+ In particular. in the lowest density bin the contribution of groups is almost negligible. while in the highest bin single and binary galaxies are almost absent.," In particular, in the lowest density bin the contribution of groups is almost negligible, while in the highest bin single and binary galaxies are almost absent."
+ Each environment. however. spans at least three of our density bins.," Each environment, however, spans at least three of our density bins."
+ In Figure 2. we show the mass functions of galaxies in different density bins. compared two by two.," In Figure \ref{pm2gc} we show the mass functions of galaxies in different density bins, compared two by two."
+ We find that the mass function depends on local density: lower density bins have proportionally a larger population of low mass galaxies than higher density regions., We find that the mass function depends on local density: lower density bins have proportionally a larger population of low mass galaxies than higher density regions.
+ The K-S test can reject the null hypothesis that the distributions are drawn from the same parent distribution when we compare DI with all the other bins. while it is inconclusive in all other cases.," The K-S test can reject the null hypothesis that the distributions are drawn from the same parent distribution when we compare D1 with all the other bins, while it is inconclusive in all other cases."
+" However. looking at the figure. it clearly emerges that 1e Slope of the D4 mass function at masses above M,/AM..~1011 is much shallower than the slope in any other density bin. that with 18 normalization adopted it is equivalent to say that in the highest density bin D4 there is an “excess” of high mass galaxies. compared o the other bins."," However, looking at the figure, it clearly emerges that the slope of the D4 mass function at masses above $M_{\ast}/M_{\odot}\sim
+10^{11}$ is much shallower than the slope in any other density bin, that with the normalization adopted it is equivalent to say that in the highest density bin D4 there is an “excess” of high mass galaxies, compared to the other bins."
+ To substantiate this on statistical grounds. since galaxies in the lowest mass bins are very numerous and they orobably strongly influence the K-S test results. we try pushing up the mass limit so to exclude those galaxies from the analysis.," To substantiate this on statistical grounds, since galaxies in the lowest mass bins are very numerous and they probably strongly influence the K-S test results, we try pushing up the mass limit so to exclude those galaxies from the analysis."
+ Redetining the mass limit entails a slightly change in the limits of he local density bins. so we compute them For logM./M.2:10.5. the differences in the mass function between D4 and the other density bins become statistically significant.," Redefining the mass limit entails a slightly change in the limits of the local density bins, so we compute them For $ \log M_{\ast}/M_{\odot}\geq
+10.5$, the differences in the mass function between D4 and the other density bins become statistically significant."
+ In general. even for a very high mass threshold (logM; 10.8). the differences in the mass functions of galaxies in different density bins remain statistically significant. showing that local density matters for any mass limit adopted.," In general, even for a very high mass threshold $ \log M_{\ast}/M_{\odot}\geq
+10.8$ ), the differences in the mass functions of galaxies in different density bins remain statistically significant, showing that local density matters for any mass limit adopted."
+ As seen in refpm2gesisto. galaxiesingroups. binarysvstemsandsinglegalaxiescov erdidd.very SCRwide.GL ering eterosenenisalmos," As seen in \\ref{pm2gc_histo}, galaxies in groups, binary systems and single galaxies cover different ranges of local densities, therefore we now wish to test whether our local density results are driven by galaxies in specific global environments (for example, only in massive groups)."
+"t4 ορύ(ρίno""etfid HRY HH gah πο μα αυ. and eui.ο 5008/5) (plots not shown)."," So we tried excluding single galaxies, or galaxies located in massive groups $\sigma_{group} >400 km/s$ and $\sigma_{group} >500 km/s$ ) (plots not shown)."
+ In all these cases we always find a similar dependence of the mass functions on the local density as we see in the general field., In all these cases we always find a similar dependence of the mass functions on the local density as we see in the general field.
+" Therefore. —qe variations of the mass distributions with local density are not ""riven by a different dependence in a specitic global environment."," Therefore, the variations of the mass distributions with local density are not driven by a different dependence in a specific global environment."
+ We use the WINGS dataset to study galaxies in clusters in the local Universe., We use the WINGS dataset to study galaxies in clusters in the local Universe.
+ Figure 3. shows the distribution of local densities in this sample and the 4 density bins., Figure \ref{wings_histo} shows the distribution of local densities in this sample and the 4 density bins.
+ In clusters. galaxies cover a range of local density of about 1.8 dex.," In clusters, galaxies cover a range of local density of about 1.8 dex."
+ In Figure 4+ we show the mass function of galaxies in the different bins of local density. compared two by Also in the case of clusters. there is a dependence of the mass function on the local density.," In Figure \ref{wings} we show the mass function of galaxies in the different bins of local density, compared two by Also in the case of clusters, there is a dependence of the mass function on the local density."
+ In. general (except when we compare D3 and D4 whose shapes are very similar). lower density bins have proportionally a greater number of lower mass galaxies.," In general (except when we compare D3 and D4 whose shapes are very similar), lower density bins have proportionally a greater number of lower mass galaxies."
+ The K-S test can reject the hypothesis of a common parent distribution with a high level of significance in most ofthe cases., The K-S test can reject the hypothesis of a common parent distribution with a high level of significance in most of the cases.
+ We note that. unlike the mass function in the general field at the same redshift. in clusters the mass functions in the highest density bin (D4. and perhaps D3) flattens out at low galaxy masses. below log/M.~10-75.," We note that, unlike the mass function in the general field at the same redshift, in clusters the mass functions in the highest density bin (D4, and perhaps D3) flattens out at low galaxy masses, below $\log M_{\ast}/M_{\odot}\sim 10.5$."
+" This is suggestive of a sort of ""deficit"" of low-massM. galaxies with respect to intermediate masses compared o lower density regions.", This is suggestive of a sort of “deficit” of low-mass galaxies with respect to intermediate masses compared to lower density regions.
+ As before. the K-S test is particularly sensitive to the large number of low mass galaxies. so we push up the mass limit in order 0 detect possible differences in the slope of the mass functions at high mass.," As before, the K-S test is particularly sensitive to the large number of low mass galaxies, so we push up the mass limit in order to detect possible differences in the slope of the mass functions at high mass."
+" After redetining the density bins. we find that local density effects in low-z clusters are not visible at intermediate-igh galaxy masses. as the K-S finds differences between the mass unction of galaxies in different density regions only for a mass imit logM,/M.|."," After redefining the density bins, we find that local density effects in low-z clusters are not visible at intermediate-high galaxy masses, as the K-S finds differences between the mass function of galaxies in different density regions only for a mass limit $\log
+M_{\ast}/M_{\odot}\leq 10.1$."
+ This limit is even lower than the mass completeness limit of the PM2GC survey., This limit is even lower than the mass completeness limit of the PM2GC survey.
+ In contrast. as we have seen in $33.1 for the PM2GC. the local density effects in the general field on the shape of the mass function do not disappear at any mass.," In contrast, as we have seen in 3.1 for the PM2GC, the local density effects in the general field on the shape of the mass function do not disappear at any mass."
+ We use the field sample of the ICBS dataset to characterize galaxies in the field at intermediate redshifts., We use the field sample of the ICBS dataset to characterize galaxies in the field at intermediate redshifts.
+ Figure 5. shows the distribution of the local density in this sample and the + density bins., Figure \ref{icbs_histo} shows the distribution of the local density in this sample and the 4 density bins.
+ In the histogram. very isolated galaxies without an estimates of local density (see refic)) are assigned οσο)—-1.5.," In the histogram, very isolated galaxies without an estimates of local density (see \\ref{ic}) ) are assigned $\log(LD)= -1.5$."
+ We ean immediately see that. excluding very isolated galaxies. for which we do not have a real estimates of local density. the range of local densities spanned is cov dex.," We can immediately see that, excluding very isolated galaxies, for which we do not have a real estimates of local density, the range of local densities spanned is very wide, covering almost 4 dex."
+ range is also very similar to ved ορ it fieldis a very environment. with very sparse regions but also with highly populated ones.," This range is also very similar to that we found for the PM2GC, indicating that actually the (general) field is a very heterogeneous environment, with very sparse regions but also with highly populated ones."
+ In Figure 6 we present the mass functions of galaxies in different density bins. compared two by two.," In Figure \ref{icbs} we present the mass functions of galaxies in different density bins, compared two by two."
+ Again. the mass function depends on local density in the sense that lower density regions have proportionally a larger population of low mass galaxies than higher density regions.," Again, the mass function depends on local density in the sense that lower density regions have proportionally a larger population of low mass galaxies than higher density regions."
+ Despite the quite small number statistic. the K-S test can always reject the null hypothesis that the distributions are drawn rom the same parent distribution except when we compare DI and D2.," Despite the quite small number statistic, the K-S test can always reject the null hypothesis that the distributions are drawn from the same parent distribution except when we compare D1 and D2."
+ Moreover. looking at the figure. as in the PM2GC. we tind tha he slope of the D4 (and maybe D3) mass function at masses above M./M.~10''* is shallower than the slope in the other density bins. indicating a possible “excess” of high mass galaxies in tha bin. compared to the other bins.," Moreover, looking at the figure, as in the PM2GC, we find that the slope of the D4 (and maybe D3) mass function at masses above $M_{\ast}/M_{\odot}\sim
+10^{11.2}$ is shallower than the slope in the other density bins, indicating a possible “excess” of high mass galaxies in that bin, compared to the other bins."
+ In this ease we decide not to further push up the mass limit. both because it is already fairly high and because the statistica uncertainty would be too large.," In this case we decide not to further push up the mass limit, both because it is already fairly high and because the statistical uncertainty would be too large."
+are differences in philosophy: for cxample. 2? obtained = TOOOKX (and = 0.9) for AW 442. and on that basis reclassified the star fromg (X3 to AT. implicitly adopting a mietallicitv-independent relationship between spectral type and temperature.,"are differences in philosophy; for example, \citet{venn99}
+ obtained = 7900K (and = 0.9) for AV 442, and on that basis reclassified the star from A3 to A7, implicitly adopting a metallicity-independent relationship between spectral type and temperature."
+ Our approach is to adopt spectral types consistently based: solely on spectral. morphology. but. to accept a metallicitv-dependent relationship between spectral tvpe and physical parameters.," Our approach is to adopt spectral types consistently based solely on spectral morphology, but to accept a metallicity-dependent relationship between spectral type and physical parameters."
+ Phat is. we would. retain the AS classification lor AV 442. but would assign νι= SOOO Ix (Table 5)).," That is, we would retain the A3 classification for AV 442, but would assign $T_{\rm eff} =$ 8000 K (Table \ref{scheme}) )."
+ Three stars in. Venn’s SAIC sample appear to have temperatures that are inconsistent with those in Table 5. given their spectral vpes. namely AY 136 (Sk 54. AO La). 478 (Sk 154. AO Lh) and 213 (Sk 75. A2 lab).," Three stars in Venn's SMC sample appear to have temperatures that are inconsistent with those in Table \ref{scheme} given their spectral types, namely AV 136 (Sk 54, A0 Ia), 478 (Sk 154, A0 Ib) and 213 (Sk 75, A2 Iab)."
+ These discrepancies appear to argelv be a result. of dillerent. published. classifications. e.g. Venn's temperatures are consistent with the spectral tvxs of 7? who gave Classifications for AY 478 (X5 lab) and AV 213 (A23 1).," These discrepancies appear to largely be a result of different published classifications, e.g. Venn's temperatures are consistent with the spectral types of \citet{hkg91} who gave classifications for AV 478 (A5 Iab) and AV 213 (A2–3 I)."
+ Figure 12. compares the solar-abundanee temperature scale cerived here with those compile by ? and by ?.., Figure \ref{scales} compares the solar-abundance temperature scale derived here with those compiled by \citet{sk82} and by \citet{hm84}.
+ Phe main cilferences are at the very earjest ancl latest types considered. here., The main differences are at the very earliest and latest types considered here.
+ The higher values we find for the late D-vpes are consistent with newer calibrations (e.g...2). so his is not a cause for concern.," The higher values we find for the late B-types are consistent with newer calibrations \citep[e.g.,][]{pc98}, so this is not a cause for concern."
+ The dillerences in the later vpes. notably F5. could be attributable to the cdilliculties of classification in this region or to refinements in the moclel atmosphere codes (for example. the inclusion of molecular ines).," The differences in the later types, notably F5, could be attributable to the difficulties of classification in this region or to refinements in the model atmosphere codes (for example, the inclusion of molecular lines)."
+ For simplicity we assumed a single relative abundance or all metallic species (ie. Q.17Z. Y: however. detailed analyses such as those of ?. find à spread of abundances for different species.," For simplicity we assumed a single relative abundance for all metallic species (i.e. $Z_\odot$ ); however, detailed analyses such as those of \citet{venn99} find a spread of abundances for different species."
+ Venn specifically highlighted. the relative uncerabundance of the alpha elements. including caleium (in contrast to the SAIC Η region. results. from ? which found no calcium. underabundance).," Venn specifically highlighted the relative underabundance of the alpha elements, including calcium (in contrast to the SMC H region results from \citet{r92} which found no calcium underabundance)."
+ To investigate the significance of adopting a lower calcium. abundance. model spectra (Z = 0.10Z .) were caleulated for 7500. 5000 and 8500 Ix. the domain in which the A line is most sensitive to temperature ancl abundance. effects.," To investigate the significance of adopting a lower calcium abundance, model spectra $Z$ = $Z_\odot$ ) were calculated for 7500, 8000 and 8500 K, the domain in which the $K$ line is most sensitive to temperature and abundance effects."
+ Such a reduction enhances the temperature effect. described. here. but. the," Such a reduction enhances the temperature effect described here, but the"
+to be at least as large as (he gas pressure μμ]. where Ap is Doltzmann's constant. (o maintain a coronal loop.,"to be at least as large as the gas pressure $2 n k_{B} T$, where $k_{B}$ is Boltzmann's constant, to maintain a coronal loop."
+" As a test of the model. Mullanetal.(2006). surveved the literature for independent measurements of T.n.L. and D,,;, for the stars in their studs. finding (hat 178 of 212 measurements were consistent wilh the IIaisch moclel."," As a test of the model, \citet{Mullan06}
+ surveyed the literature for independent measurements of $T, n, L$, and $B_{min}$ for the stars in their study, finding that 178 of 212 measurements were consistent with the Haisch model."
+ This justifies treating the derived parameters as valid independent of the IHaisch model., This justifies treating the derived parameters as valid independent of the Haisch model.
+" We first verify that the IHaisch model gives reasonable restills for solar parameters,", We first verify that the Haisch model gives reasonable results for solar parameters.
+ Eruptive solar [lares have 7;104? sec and LAL~107en? (Priest&Forbes2000)., Eruptive solar flares have $\tau_{d} \sim 10^{4-5}$ sec and $EM \sim 10^{49-50} {\rm cm}^{-3}$ \citep{Priest00}.
+. Using these values. the Iaisch model predicts post-fDare parameters of 7~(4—13) ALN. nc(0.3—5.6)x10! em.7. Le10!H. em. and Bain(10—70) G. Compact solar flares have 7j107 sec and EAL~107.IP em7? (Priest&Forbes 2000)..," Using these values, the Haisch model predicts post-flare parameters of $T
+\sim (4-13)$ MK, $n \sim (0.3-5.6) \times 10^{10}$ ${\rm cm}^{-3}$, $L
+\sim 10^{10-11}$ cm, and $B_{min} \sim (10 - 70)$ G. Compact solar flares have $\tau_{d} \sim 10^{3}$ sec and $EM \sim 10^{47-49}$ ${\rm
+cm}^{-3}$ \citep{Priest00}. ."
+" Using these values. the llaisch ""magives T—(1413) MK. ne(3—5)x10H em7. E—(0.51.6)x10? em. and D,,5,~(90—220) G. These ranges of T.n.L. and D,,;, are consistent with independent empirical values obtained from images and. X-ray. cata for Leaving loops in the sun Yokovama2002).."," Using these values, the Haisch model gives $T \sim (4-13)$ MK, $n \sim (3 - 5) \times 10^{11}$ ${\rm
+cm}^{-3}$, $L \sim (0.5 - 1.6) \times 10^{9}$ cm, and $B_{min} \sim
+(90 - 220)$ G. These ranges of $T, n, L,$ and $B_{min}$ are consistent with independent empirical values obtained from images and X-ray data for flaring loops in the sun \citep{Feldman95,Shibata02}. ."
+ To calculate 95e and d; from the Iaisch model. we use {μμ lor the upstream magnetic field and 6» for the density.," To calculate $\delta_{SP}$ and $d_{i}$ from the Haisch model, we use $B_{min}$ for the upstream magnetic field and $n$ for the density."
+ Sweet-Parker current sheets extend {ο system scales 1986).. so Lap is on the order of the coronal loop radius AM?~L/10. consistent with the llaisch model.," Sweet-Parker current sheets extend to system scales \citep{Biskamp86}, so $L_{SP}$ is on the order of the coronal loop radius $A^{1/2} \sim L
+/ 10$, consistent with the Haisch model."
+" Finally. we use 7 to caleulate the Spitzer resistivity (Spitzer&ILirm1953) where m, is the electron mass and InA=In((3/26(5Tπλ is the Coulomb logarithm."," Finally, we use $T$ to calculate the Spitzer resistivity \citep{Spitzer53}
+ where $m_{e}$ is the electron mass and $\ln \Lambda = \ln [ (3 / 2
+e^{3}) (k_{B}^{3} T^{3} / \pi n)^{1/2}]$ is the Coulomb logarithm."
+" Use of this formula is justified because the electron mean [ree path. km lor solar conditions. where cj, isthe electron thermal speed aud 7; is the elec(von-ion collision frequency) is small compared to length scales in the outflow direction (Lap10 km) and along the current sheet (L—10? km)."," Use of this formula is justified because the electron mean free path $\lambda_{{\rm mfp,e}} \sim v_{th,e} / \nu_{ei} \sim 25$ km for solar conditions, where $v_{th,e}$ isthe electron thermal speed and $\nu_{ei}$ is the electron-ion collision frequency) is small compared to length scales in the outflow direction $L_{SP} \sim 10^{4}$ km) and along the current sheet $L ~ \sim 10^{5}$ km)."
+ The result of comparing dsp to d; using the stellar [Iare is plotted in Fig. 3.., The result of comparing $\delta_{SP}$ to $d_{i}$ using the stellar flare data is plotted in Fig. \ref{spvsdi}.
+ Representative solar values based onty=1057? sec and EA=10/?cmCla* for eruptive flares (OgpcLLO em and d;e200 em) and 7;=10* sec and EM=I0Pem? for compact flares (Osp~44 cm and d;~352 em) are plotted as the asterisk and plus. respectively.," Representative solar values based on $\tau_{d} = 10^{4.5}$ sec and $EM = 10^{49.5} {\rm
+cm}^{-3}$ for eruptive flares $\delta_{SP} \sim 110$ cm and $d_{i}
+\sim 200$ cm) and $\tau_{d} = 10^{3}$ sec and $EM = 10^{48} {\rm
+cm}^{-3}$ for compact flares $\delta_{SP} \sim 44$ cm and $d_{i} \sim
+35$ cm) are plotted as the asterisk and plus, respectively."
+ A dashed line with slope of unity is plotted., A dashed line with slope of unity is plotted.
+ The agreement is extremely good., The agreement is extremely good.
+ A least squares analysis gives a best fit slope of 0.98+0.02 with a correlation coefficient of 0.981., A least squares analysis gives a best fit slope of $0.98 \pm 0.02$ with a correlation coefficient of 0.981.
+ It is encouraging that the slope of the line in Fig., It is encouraging that the slope of the line in Fig.
+ 3. is consistent with unity., \ref{spvsdi} is consistent with unity.
+ However. there are ambiguitües in the dala analysis.," However, there are ambiguities in the data analysis."
+" For example. weused d; as the critical length scale. whereas p, is more applicable to the corona (but more difficult to estimate)."," For example, weused $d_{i}$ as the critical length scale, whereas $\rho_{s}$ is more applicable to the corona (but more difficult to estimate)."
+" These scales diller⋅⊳ by a factor. of2e sb 9,77. where 2j, is. the ratio. of. gas pressure to totalmagnetic"," These scales differ by a factor of $\beta_{tot}^{1/2}$ , where $\beta_{tot}$ is the ratio of gas pressure to totalmagnetic"
+" These scales diller⋅⊳ by a factor. of2e sb 9,77. where 2j, is. the ratio. of. gas pressure to totalmagnetic."," These scales differ by a factor of $\beta_{tot}^{1/2}$ , where $\beta_{tot}$ is the ratio of gas pressure to totalmagnetic"
+The template-fitting photometric redshift technique makes use of the available ancl reasonably detailed knowledge of galaxy SEDs and in principle it may be used reliably even for populations of galaxies for which there are [ow or no spectroscopically confirmed: redshifts.,The template-fitting photometric redshift technique makes use of the available and reasonably detailed knowledge of galaxy SEDs and in principle it may be used reliably even for populations of galaxies for which there are few or no spectroscopically confirmed redshifts.
+ However. crucial to its success. is the compilation of a library of accurate and representative template SEDs (see HlIloge et 11998: Firth 2002b).," However, crucial to its success, is the compilation of a library of accurate and representative template SEDs (see Hogg et 1998; Firth 2002b)."
+ Enipirical templates are typically derived from nearby bright galaxies. which may not be truly representative of high redshift. galaxies.," Empirical templates are typically derived from nearby bright galaxies, which may not be truly representative of high redshift galaxies."
+ Conversely. while theoretical SEDs can cover a large range of star formation histories. metallicities. dust extinction. models etc..," Conversely, while theoretical SEDs can cover a large range of star formation histories, metallicities, dust extinction models etc.,"
+. not all combinations of these parameters (at any particular redshift) are realistic. and the ad hoc inclusion of superfluous templates increases the potential for misidentifications when using observations with noisy. photometry.," not all combinations of these parameters (at any particular redshift) are realistic, and the ad hoc inclusion of superfluous templates increases the potential for misidentifications when using observations with noisy photometry."
+ An alternative approach can be used when one has a sullicienthy large 72100.1000. depending on the redshift range) and representative subsample with spectroscopic redshifts.," An alternative approach can be used when one has a sufficiently large $\sim$ 100–1000, depending on the redshift range) and representative subsample with spectroscopic redshifts."
+ Then one can fit a polvnomial or other function mapping the photometric data to the known redshifts and use this to estimate redshifts for the remainder of the sample with unknown redshifts. CConnolly et. 11995b: Munner. Szalav Connolly 2000: Sowards-Enmumoerd ct 22000).," Then one can fit a polynomial or other function mapping the photometric data to the known redshifts and use this to estimate redshifts for the remainder of the sample with unknown redshifts Connolly et 1995b; Brunner, Szalay Connolly 2000; Sowards-Emmerd et 2000)."
+ With this approach. errors in the estimated redshifts mav also be estimated. analytically or via Monte. Carlo simulations.," With this approach, errors in the estimated redshifts may also be estimated analytically or via Monte Carlo simulations."
+ An extension of the latter approach is to use Artificia Neural Networks (ANNs hereafter)., An extension of the latter approach is to use Artificial Neural Networks (ANNs hereafter).
+ ANNs have been usec before in astronomy for. amongst other things. galaxy morphological classification. SStorrie-Lombardi et 11992: Naim et 11995: Lahav et 11996). morphologica starfealaxy separation BBertin Arnouts 1996: Anedveon et 22000) and stellar spectral classification BBailer-Jones. Irwin von Llippel 1998: Allende Prieto e 22000: Weaver 2000).," ANNs have been used before in astronomy for, amongst other things, galaxy morphological classification Storrie-Lombardi et 1992; Naim et 1995; Lahav et 1996), morphological star/galaxy separation Bertin Arnouts 1996; Andreon et 2000) and stellar spectral classification Bailer-Jones, Irwin von Hippel 1998; Allende Prieto et 2000; Weaver 2000)."
+ Essentially an ANN takes a se of inputs logarithms of Ηχος mmaegnitucdes in dilferent filters) for each object. applies some non-linear function. and outputs a value tthe estimated redshift).," Essentially an ANN takes a set of inputs logarithms of fluxes – magnitudes – in different filters) for each object, applies some non-linear function, and outputs a value the estimated redshift)."
+ The ANN is first trained tthe coefficients (weights) of he function are optimized. by using a training set. where he desired. output is known., The ANN is first trained – the coefficients (weights) of the function are optimized – by using a training set where the desired output is known.
+ The ANN may then be used on any number of other objects with similar inputs nunagnituces in the same filter set) but unknown outputs rredshifts)., The ANN may then be used on any number of other objects with similar inputs magnitudes in the same filter set) but unknown outputs redshifts).
+ As well as using all of the information contained in he magnitudes ancl colours. provided. the training set is a representative subsample of the data. the ANN will also ake into account the Bavesian priors on the galaxy recishift distribution BBenitez 1998: Teplitz et 22001).," As well as using all of the information contained in the magnitudes and colours, provided the training set is a representative subsample of the data, the ANN will also take into account the Bayesian priors on the galaxy redshift distribution Benitez 1998; Teplitz et 2001)."
+ While choosing a template library that is both sullicient ancl non-superfluous is a source of concern for the template-fitting method. ANNs automatically fit the true range of galaxy SEDs.," While choosing a template library that is both sufficient and non-superfluous is a source of concern for the template-fitting method, ANNs automatically fit the true range of galaxy SEDs."
+ Another potential advantage of ANNs relative to the template-fitting method is that the weights applied to each filter⋅ may be more optimal. than simple. 2 \7-weighting., Another potential advantage of ANNs relative to the template-fitting method is that the weights applied to each filter may be more optimal than simple $\chi^2$ -weighting.
+. In addition one can also feed in other observational input such as image size or surface brightness. morphology. anc concentration parameters where such data are available.," In addition one can also feed in other observational input such as image size or surface brightness, morphology and concentration parameters where such data are available."
+ E is interesting then to see how the two methods compare., It is interesting then to see how the two methods compare.
+ This paper explores the use of ANNs as a potentia tool for photometric redshift’ determination., This paper explores the use of ANNs as a potential tool for photometric redshift determination.
+ Ehe lavout of this paper is as follows., The layout of this paper is as follows.
+ In ὃν the ANNs are. describe and in ὃς a semi-analvtic model (used το provide a simulated. galaxy catalogue) is introduced., In $\S$ \ref{sec.anns} the ANNs are described and in $\S$ \ref{sec.sam} a semi-analytic model (used to provide a simulated galaxy catalogue) is introduced.
+ In 54 — the ANN parameters. (architecture ancl training set size) are investigated using the simulated galaxy catalogue and in 85 the performance of ANNs are compared. with the performance of the traditional template-itting method., In $\S$ \ref{sec.cf} the ANN parameters (architecture and training set size) are investigated using the simulated galaxy catalogue and in $\S$ \ref{sec.hyperz} the performance of ANNs are compared with the performance of the traditional template-fitting method.
+ 56 looks at the elfect of photometric noise and in ὃν ANNs are investigated as a method for also determining spectral tvpe from redshifted data., $\S$ \ref{sec.hist} looks at the effect of photometric noise and in $\S$ \ref{sec.sed} ANNs are investigated as a method for also determining spectral type from redshifted data.
+ En 88.. ANNs are tested on Sloan Digital Sky Survey observational data.," In $\S$ \ref{sec.sdss}, ANNs are tested on Sloan Digital Sky Survey observational data."
+ Phe science prospects are briefly discussed in $9.., The science prospects are briefly discussed in $\S$ \ref{sec.conclude}.
+ An ANN comprises a set of input nodes. one or more output nodes. and one or more hidden lavers each containing a number of nodes reffig.photz.arch:: see BBishop 1995. and references therein. for background).," An ANN comprises a set of input nodes, one or more output nodes, and one or more hidden layers each containing a number of nodes \\ref{fig.photz.arch}; see Bishop 1995, and references therein, for background)."
+" A particular network architecture may be clenotec by Αλ Nerina where Ni, ds the number of input nodes. Ny is the number of nodes in the first. hidden laver. and so on."," A particular network architecture may be denoted by $N_{\mathrm{in}}$ $N_1$ $N_2$ $N_{\mathrm{out}}$ where $N_{\mathrm{in}}$ is the number of input nodes, $N_1$ is the number of nodes in the first hidden layer, and so on."
+ For example 9:6:1. takes 9 inputs. has 6 nodes in a single hidden laver and gives a single output.," For example 9:6:1 takes 9 inputs, has 6 nodes in a single hidden layer and gives a single output."
+ The nodes are connected and each connection carries a weight which together comprise the vector of coellicients w which are to be optimized., The nodes are connected and each connection carries a weight which together comprise the vector of coefficients $\bmath{w}$ which are to be optimized.
+ Unless otherwise stated. here," Unless otherwise stated, here"
+Science Foundation through TeraCirid resources provided by NICS Wraken under grant nuuber TO-ASTIOO00LO.,Science Foundation through TeraGrid resources provided by NICS Kraken under grant number TG-AST100040.
+ AS is supported by NSF evant AST-0S07381 aud NASA erauts NNNOQATO5C. and NNXIOÀQ039€.. We. thas the referee. Ioannis C'ontopoulos. for comunents that helped prove this paper.," AS is supported by NSF grant AST-0807381 and NASA grants NNX09AT95G and NNX10A039G. We thank the referee, Ioannis Contopoulos, for comments that helped improve this paper."
+ The simulations preseuted iu this paper were performed ou computational resources supported bv the PICSGE-OIT Tiel Performance Computing Center aud Visualization Laboratory., The simulations presented in this paper were performed on computational resources supported by the PICSciE-OIT High Performance Computing Center and Visualization Laboratory.
+ This research used resources of the National Energy Rescarcli Scientific Computing Center. which is supported by the Office of Science of the US Department of Energy uudoer contract No.," This research used resources of the National Energy Research Scientific Computing Center, which is supported by the Office of Science of the US Department of Energy under contract No."
+ DE-AC02-05CTH1231., DE-AC02-05CH11231.
+ We derive here in full generality our resistive current prescription., We derive here in full generality our resistive current prescription.
+ We start in the fluid frame. with charge py aud curent flowine along the electric field with magnitude oZ.," We start in the fluid frame, with charge $\rho_0$ and current flowing along the electric field with magnitude $\sigma E_0$."
+ For convenience. we pick the current aud electric field to point alone the positive z axis.," For convenience, we pick the current and electric field to point along the positive z axis."
+ The maeguetic field can point along the positive or negative z axis. depending ou whether By is positive or negative.," The magnetic field can point along the positive or negative z axis, depending on whether $B_0$ is positive or negative."
+ If we boost along the z axis with 4=(0.0.j02). the current Lb vector in the boosted frame satisfies Boosting again iu the x direction transverse to the electric and maguetie fields with ο=(2.0.0). we obtain for lab frame quantities the svstemi of equations," If we boost along the z axis with $\vec{\beta_1}=(0,0,\beta_z)$, the current 4 vector in the boosted frame satisfies Boosting again in the x direction transverse to the electric and magnetic fields with $\vec{\beta_2}=(\beta_x,0,0)$, we obtain for lab frame quantities the system of equations"
+range involves electron energies lower than those involved in producing observable radio svachrotron emission of the lobes.,range involves electron energies lower than those involved in producing observable radio synchrotron emission of the lobes.
+ Searching for analogous relict Ceghiost) structures in other evolved radio galaxies and radio-loud quasars using current X-ray instruments has been in fact alreacly proposed. and considered in the literatureFabian2011:Moczetal. 2011).," Searching for analogous relict (`ghost') structures in other evolved radio galaxies and radio-loud quasars using current X-ray instruments has been in fact already proposed and considered in the literature\citep[e.g.,][]{fab09,blu11,moc11}."
+. In ihe previous section 33.3 we have noted a relatively voung age of the observed radio structure of emerging from the model fits. /2z35 MAIve. which is in between of the ages evaluated for the inner and outer lobes of the analvzed DDRGs.," In the previous section 3.3 we have noted a relatively young age of the observed radio structure of emerging from the model fits, $t \simeq 35$ Myr, which is in between of the ages evaluated for the inner and outer lobes of the analyzed DDRGs."
+ Correspondinely. a relatively high expansion velocily (feeD/(2e1)~0.2 has been found for the source. about ten times higher that the analogous values estimated for the outer lobes in the control sample of DDRGs. and a factor of at least a few higher (han (the expansion velocity of the inner lobes in such svstens.," Correspondingly, a relatively high expansion velocity $v_{\rm h}(t)/c \simeq D/(2 c\,t) \simeq 0.2$ has been found for the source, about ten times higher that the analogous values estimated for the outer lobes in the control sample of DDRGs, and a factor of at least a few higher than the expansion velocity of the inner lobes in such systems."
+ Are however these somehow extreme values realistic at all. or are thev just artifacts of the incorrect scenario applied?," Are however these somehow extreme values realistic at all, or are they just artifacts of the incorrect scenario applied?"
+ In all (he kinematic models for LILI-Utvpe sources. characteristic asvimnmetries in lengths and luminosities between the (wo opposite lobes are expected due to the light (me travel ellects (e.g..Longair&Rilev1979:Bestetal.1995:ScheuerRvs2000).," In all the kinematic models for II-type sources, characteristic asymmetries in lengths and luminosities between the two opposite lobes are expected due to the light time travel effects \citep[e.g.,][]{lon79,bes95,sch95,rys00}."
+. The size ratio between the longer (jet side) and shorter (counter-]et. side) lobes is in particular (oy/06)cos 0]. giving the advance speed of the jets In the case of we have estimated q71.078. leading to ey20.0315e/cos0 and thus implving 90< 88.," The size ratio between the longer (jet side) and shorter (counter-jet side) lobes is in particular $q=[1+(v_{\rm h}/c) \cos\theta]/[1-(v_{\rm h}/c) \cos\theta]$ , giving the advance speed of the jets In the case of we have estimated $q \simeq 1.078$, leading to $v_{\rm h} \simeq 0.0375 \, c /\cos\theta$ and thus implying $\theta<88\degr$ ."
+ Η we further assume lobes. inclination in the source 8~80°. the expansion velocity reads as ej~0.2166. and hence the source age as MMvr. ," If we further assume lobes' inclination in the source $\theta\sim 80\degr$, the expansion velocity reads as $v_{\rm h}\sim 0.216\,c$, and hence the source age as $t \simeq D / (2\,v_{\rm h}) \sim 35.4$ Myr."
+Thus. the kinematic estimates of the age and of the advance velocity of the observed radio structure in are filly consistent with the values resulting from the dynamical modeling presented in 23.3.," Thus, the kinematic estimates of the age and of the advance velocity of the observed radio structure in are fully consistent with the values resulting from the dynamical modeling presented in 3.3."
+ We note that the jet inclination 080° is in [act more probable than 990° adopted in (he fitting procedure also for the other DDRGs analvzed., We note that the jet inclination $\theta \simeq 80\degr$ is in fact more probable than $\theta \simeq 90\degr$ adopted in the fitting procedure also for the other DDRGs analyzed.
+ However. the resulting model parameters are not sensitive to suchrelatively minor changes in (he assumed value of 0. since," However, the resulting model parameters are not sensitive to suchrelatively minor changes in the assumed value of $\theta$ , since"
+parameter of B/n'4=2 uG en. solar abundance.,"parameter of $B/n^{1/2}=2$ $\mu$ G $\mathrm{cm}^{3/2}$, solar abundance."
+ As shown in Figure 4. both pure shock models and shock plus precursor models can not explain the observed distribution of line ratios. at variance with pure photoionization models.," As shown in Figure 4, both pure shock models and shock plus precursor models can not explain the observed distribution of line ratios, at variance with pure photoionization models."
+ We therefore conclude that the dominant tonization mechanism of gas clouds emittingiv.Her. and Cit] is photoionization.," We therefore conclude that the dominant ionization mechanism of gas clouds emitting, and ] is photoionization."
+ To investigate the relationship between metallicity and redshift or lummosity. we divided our sample into the following three redshift bins: 1.2<i<2.0. 2.0<22.5. and 2.5<cz3.8.," To investigate the relationship between metallicity and redshift or luminosity, we divided our sample into the following three redshift bins: $1.2 < z < 2.0$, $2.0 < z < 2.5$, and $2.5 < z < 3.8$."
+ As for the luminosity. since the non-thermal AGN continuum is obscured in HzRGs. we used the line luminosity as a proxy for the AGN’s bolometric luminosity.," As for the luminosity, since the non-thermal AGN continuum is obscured in HzRGs, we used the line luminosity as a proxy for the AGN's bolometric luminosity."
+ This directly traces the rate of He -10nizing photons which is proportional to the bolometric luminosity., This directly traces the rate of $\mathrm{He}^+$ -ionizing photons which is proportional to the bolometric luminosity.
+ Luminosity bins are as follows: 41.5«log L(Hei) <42.5. 42.5<log L(Hen) <43.0 and. 43.0<log (Hem) <45.0 (see. Table 3).," Luminosity bins are as follows: $41.5 < \log L$ ) $< 42.5$, $42.5 < \log L$ ) $< 43.0$ and, $43.0 < \log L$ ) $< 45.0$ (see, Table 3)."
+ In Figures 5 and 6. we investigate the averaged flux ratios of and for each luminosity and redshift bin to examine the possible metallicity dependence on luminosity and redshift.," In Figures 5 and 6, we investigate the averaged flux ratios of and for each luminosity and redshift bin to examine the possible metallicity dependence on luminosity and redshift."
+ There ts 10 clear dependence betwee the flux ratios and redshift as shown in Figure 5. suggesting that the NLRs metallicity of HzRGs does not show significant redshift evolution in the range of |€z4.," There is no clear dependence between the flux ratios and redshift as shown in Figure 5, suggesting that the NLRs metallicity of HzRGs does not show significant redshift evolution in the range of $1 \la z \la 4$."
+ On the other hand. in Figure 6 we find that the flux ratios are correlated with the luminosity. suggesting that more luminous HzRGs have more metal-rich gas clouds.," On the other hand, in Figure 6 we find that the flux ratios are correlated with the luminosity, suggesting that more luminous HzRGs have more metal-rich gas clouds."
+ Figure 7 shows the redshift evolution of average metallicity as derived from our sample., Figure 7 shows the redshift evolution of average metallicity as derived from our sample.
+ Here we use only the HzRG samples with log L(Hei) =42.5—43.0 to minimize possible Malmquist bias., Here we use only the HzRG samples with $\log L$ ) $=42.5-43.0$ to minimize possible Malmquist bias.
+ Note that the metallicity shown here are derived by assuming ny10Η (see Table 3)., Note that the metallicity shown here are derived by assuming $n_\mathrm{H}=10^4 cm^{-3}$ (see Table 3).
+ Figure 7 clearly shows that there is no significant decline in the NLR metallicity at high redshift., Figure 7 clearly shows that there is no significant decline in the NLR metallicity at high redshift.
+ In the same way. Figure 8 shows the luminosity dependence of the averaged NLR metallicity in our sample.," In the same way, Figure 8 shows the luminosity dependence of the averaged NLR metallicity in our sample."
+ Figure 8 clearly shows the correlation between the AGN luminosity and the NLR metallicity., Figure 8 clearly shows the correlation between the AGN luminosity and the NLR metallicity.
+ These results are completely consistent with the results of Nagao et al. (, These results are completely consistent with the results of Nagao et al. (
+2006a) but more reliable since our new sample includes a larger number of HzRGs at high redshift (7> 2.7).,2006a) but more reliable since our new sample includes a larger number of HzRGs at high redshift $z > 2.7$ ).
+ Note that the positive correlation between the AGN luminosity and the gas metallicity of ionized clouds is indicated also for the BLR. that does not show any redshift evolution (e.g.. Hamann Ferland 1993: Nagao et al.," Note that the positive correlation between the AGN luminosity and the gas metallicity of ionized clouds is indicated also for the BLR, that does not show any redshift evolution (e.g., Hamann Ferland 1993; Nagao et al."
+ 2006c: see also Jiang et al., 2006c; see also Jiang et al.
+ 2007: Juarez et al., 2007; Juarez et al.
+ 2009)., 2009).
+ In section 5.1. we reported a significant positive correlation between the NLR metallicity and the AGN luminosity.," In section 5.1, we reported a significant positive correlation between the NLR metallicity and the AGN luminosity."
+ In this section we discuss two possible origins for this correlation: We discuss each of these scenarios below., In this section we discuss two possible origins for this correlation: We discuss each of these scenarios below.
+ One possible scenario for the positive correlation between NLR metallicity and AGN luminosity is that it reflects the relation between galaxy mass and metallicity (e.g.. Lequeux et al.," One possible scenario for the positive correlation between NLR metallicity and AGN luminosity is that it reflects the relation between galaxy mass and metallicity (e.g., Lequeux et al."
+ 1979; Tremonti et al., 1979; Tremonti et al.
+ 2004: Lee et al., 2004; Lee et al.
+ 2006)., 2006).
+ Since the metal content of NLRs is the result of the past star-formation history in the host galaxies. the metallicity of galaxies and that of NLRs should be closely related.," Since the metal content of NLRs is the result of the past star-formation history in the host galaxies, the metallicity of galaxies and that of NLRs should be closely related."
+ Galaxies are characterized by a well defined mass-metallicity relation even at high redshift (Maiolino et al., Galaxies are characterized by a well defined mass-metallicity relation even at high redshift (Maiolino et al.
+ 2008)., 2008).
+ On the other hand. AGN luminosities and host galaxy masses should also be closely related. if the Eddington ratio ts roughly the same within this class of objects and the correlation between the mass of supermassive black holes (SMBHs) and that of galaxy spheroidal components holds also at high redshift.," On the other hand, AGN luminosities and host galaxy masses should also be closely related, if the Eddington ratio is roughly the same within this class of objects and the correlation between the mass of supermassive black holes (SMBHs) and that of galaxy spheroidal components holds also at high redshift."
+ Therefore. by assuming a narrow range of the Eddington ratios and the Mypg-galaxy nass relation. the positive correlation between NLR metallicity and AGN luminosity is naturally expected.," Therefore, by assuming a narrow range of the Eddington ratios and the $M_\mathrm{BH}$ -galaxy mass relation, the positive correlation between NLR metallicity and AGN luminosity is naturally expected."
+ Indeed. some observational studies report that the Eddington ratio of high-z quasars is limited in a narrow range (e.g.. Kollmeier et al.," Indeed, some observational studies report that the Eddington ratio of high-z quasars is limited in a narrow range (e.g., Kollmeier et al."
+ 2006: Trump et al., 2006; Trump et al.
+ 2009)., 2009).
+ However. this scenario has a serious problem.," However, this scenario has a serious problem."
+ We found that the NLR metallicity shows no significant redshift evolution. up to z~4.," We found that the NLR metallicity shows no significant redshift evolution, up to $z\sim4$."
+ On the contrary. the mass-metallicity relation in galaxies shows a significant redshift evolution. at least in the O<z«3 range (e.g.. Savaglio et al.," On the contrary, the mass-metallicity relation in galaxies shows a significant redshift evolution, at least in the $0 < z < 3$ range (e.g., Savaglio et al."
+ 2005: Erb et al., 2005; Erb et al.
+ 2006: Liu et al., 2006; Liu et al.
+ 2008: Maiolino et al., 2008; Maiolino et al.
+ 2008)., 2008).
+ If the relation between NLR metallicity and AGN luminosity is simply caused by the galaxy mass-metallicity relation. the NLR metallicity is also expected to show a significant redshift evolution. in contrast with our observational results.," If the relation between NLR metallicity and AGN luminosity is simply caused by the galaxy mass-metallicity relation, the NLR metallicity is also expected to show a significant redshift evolution, in contrast with our observational results."
+ Below we discuss possible reasons for this apparent contradiction and the implications of the relation between NLR metallicity and AGN luminosity., Below we discuss possible reasons for this apparent contradiction and the implications of the relation between NLR metallicity and AGN luminosity.
+ A possibility to reconcile this problem is that the apparent lack of evolution 1s the extreme consequence of galaxy downsizing on metal enrichment., A possibility to reconcile this problem is that the apparent lack of evolution is the extreme consequence of galaxy downsizing on metal enrichment.
+ We know that low mass galaxies evolve slowly. even from the chemical point of view. on a prolonged time scale extending to the current epoch.," We know that low mass galaxies evolve slowly, even from the chemical point of view, on a prolonged time scale extending to the current epoch."
+ In contrast. massive galaxies reach their chemical maturity on short time scales and at high redshift (Maiolino et al.," In contrast, massive galaxies reach their chemical maturity on short time scales and at high redshift (Maiolino et al."
+ 2008 and references therein)., 2008 and references therein).
+ Since the host galaxies of HzRGs are very massive (e.g. De Breuck et al., Since the host galaxies of HzRGs are very massive (e.g. De Breuck et al.
+ 2002). their chemical evolution may be completed at much earlier epochs than observed in the current sample. ie. at z>3.," 2002), their chemical evolution may be completed at much earlier epochs than observed in the current sample, i.e. at $z > 3$."
+ This scenario predicts that the NLR metallicities in lower-luminosity HzRGs (hence probably hosted in less massive hosts) may show evolution even at τς3., This scenario predicts that the NLR metallicities in lower-luminosity HzRGs (hence probably hosted in less massive hosts) may show evolution even at $z < 3$.
+ This prediction can be tested with sensitive spectroscopic observations of faint HzZRGs., This prediction can be tested with sensitive spectroscopic observations of faint HzRGs.
+ Alternatively. the relation between NLR metallicity and AGN luminosity may be independent of the mass-metallicity relation of galaxies. since the Eddington ratio is likely not universal.," Alternatively, the relation between NLR metallicity and AGN luminosity may be independent of the mass-metallicity relation of galaxies, since the Eddington ratio is likely not universal."
+ Recently it was reported that the BLR metallicity in quasars Is correlated with the Eddington ratio (e.g.. Shemmer et al.," Recently it was reported that the BLR metallicity in quasars is correlated with the Eddington ratio (e.g., Shemmer et al."
+ 2004; see also Nagao et al., 2004; see also Nagao et al.
+ 2002b; Shemmer Netzer 2002)., 2002b; Shemmer Netzer 2002).
+ This suggests that the NLR metallicity may be also correlated with the Eddington ratio. although the physical origin is not clear.," This suggests that the NLR metallicity may be also correlated with the Eddington ratio, although the physical origin is not clear."
+ In this case. the correlation between the," In this case, the correlation between the"
+Zaritsky et al. (,Zaritsky et al. (
+2006a. 20060. 2008) define the fundamental manifold relation as In this formulation. log(T!) is the effective mass-to-light ratio parametrized in terms of log() and log(Ll;).,"2006a, 2006b, 2008) define the fundamental manifold relation as In this formulation, $\rm log({\Upsilon}^{f}_{\rm e})$ is the effective mass-to-light ratio parametrized in terms of $\rm log(\sigma)$ and $\rm log(\rm I_e)$ ."
+ That is. logCr!)= logCr'dog(r).log(l;).," That is, $\rm
+log({\Upsilon}^{f}_{\rm e})=log({\Upsilon}^{f}_{\rm e}(\rm log(\sigma),\rm
+log(\rm I_e))$."
+ The parametrization. is determined from a fit of dynamically derived M/L ratio (using the virial theorem) as a function of log(c) and log(I;) (see Zaritsky et al., The parametrization is determined from a fit of dynamically derived M/L ratio (using the virial theorem) as a function of $\rm log(\sigma)$ and $\rm log(\rm I_e)$ (see Zaritsky et al.
+ 2008)., 2008).
+ It is clear that the exact functional shape of logCr!(log(r).log(l;)) — and hence the location of the fundamental manifold — depends on which stellar systems are included in the fit.," It is clear that the exact functional shape of $\rm log({\Upsilon}^{f}_{\rm e}(\rm log(\sigma),\rm log(\rm I_e))$ — and hence the location of the fundamental manifold — depends on which stellar systems are included in the fit."
+ For example. the original formulation of the manifold (Zaritsky et al.," For example, the original formulation of the manifold (Zaritsky et al."
+ 2006a) does not include the heavily dark matter dominated Local Group dwarf spheroidal galaxies in the fit (Zaritsky et al., 2006a) does not include the heavily dark matter dominated Local Group dwarf spheroidal galaxies in the fit (Zaritsky et al.
+ 2006b. Simon Geha 2007).," 2006b, Simon Geha 2007)."
+ A revised formulation extending to the very large M/L of the dSphs was presented in Zaritsky et al. (, A revised formulation extending to the very large M/L of the dSphs was presented in Zaritsky et al. (
+2008: see Table 1 of that paper).,2008; see Table 1 of that paper).
+ We are interested in the link between UCDs and canonical galaxies on the one hand. and on the relation between UCDs and star clusters on the other hand.," We are interested in the link between UCDs and canonical galaxies on the one hand, and on the relation between UCDs and star clusters on the other hand."
+ Therefore. we show in Fig.," Therefore, we show in Fig."
+ I4 the location of UCDs and GCs with respect to three different formulations of the fundamental manifold., \ref{manifold} the location of UCDs and GCs with respect to three different formulations of the fundamental manifold.
+ The functional form of logCr'tlogtor).log) for these three representations is indicated in Table 4..," The functional form of $\rm
+log({\Upsilon}^{f}_{\rm e}(\rm log(\sigma),\rm log(\rm I_e))$ for these three representations is indicated in Table \ref{manifold_coef}."
+" The first formulation does not include Local Group dwarf galaxies to the fit of logr'tlogtr).log(1,). nor UCDs and GCs."," The first formulation does not include Local Group dwarf galaxies to the fit of $\rm log({\Upsilon}^{f}_{\rm e}(\rm log(\sigma),\rm
+log(\rm I_e))$, nor UCDs and GCs."
+ This is the original manifold version from Zaritsky et al. (, This is the original manifold version from Zaritsky et al. (
+2006a).,2006a).
+" It is intriguing that in this formulation. by more than a decade in re. down tor,~5—7 pe (note that r, is shown in units of kpe in Fig. 14))."," It is intriguing that in this formulation, by more than a decade in $r_e$, down to $r_e \sim 5-7$ pc (note that $r_e$ is shown in units of kpc in Fig. \ref{manifold}) )."
+ Together with all other spheroids they follow a well defined linear function slightly inclined with respect to the original fundamental manifold. with a slope 0.92 + 0.01.," Together with all other spheroids they follow a well defined linear function slightly inclined with respect to the original fundamental manifold, with a slope 0.92 $\pm$ 0.01."
+ A possible interpretation of this is that for the faintest dwarf spheroidals. the dark matter halo is de-coupled from the baryons (see also Zaritsky et al.," A possible interpretation of this is that for the faintest dwarf spheroidals, the dark matter halo is de-coupled from the baryons (see also Zaritsky et al."
+ 2006b) such that the continous relation of baryon packing efficiency vs. galaxy scale breaks down., 2006b) such that the continous relation of baryon packing efficiency vs. galaxy scale breaks down.
+ Another possibility is that the faintest dwarfs are out of dynamical equilibrium., Another possibility is that the faintest dwarfs are out of dynamical equilibrium.
+ This aspect ts closely related to the discussion of the origin of dwarf satellite galaxies (dark-matter dominated cosmological substructure vs. tidal dwarf galaxy. see Kroupa et al.," This aspect is closely related to the discussion of the origin of dwarf satellite galaxies (dark-matter dominated cosmological substructure vs. tidal dwarf galaxy, see Kroupa et al."
+ 2005)., 2005).
+ The second formulation includes all objects in Fig., The second formulation includes all objects in Fig.
+ 14 for the fitting logrllog).log(l;) (Zaritsky. private communication).," \ref{manifold} for the fitting $\rm log({\Upsilon}^{f}_{\rm e}(\rm log(\sigma),\rm
+log(\rm I_e))$ (Zaritsky, private communication)."
+" Again. UCDs follow the manifold line. and only for log(r.)€—2.2 they start to ""bend down’."," Again, UCDs follow the manifold line, and only for $\rm log(r_e)\lesssim-2.2$ they start to 'bend down'."
+ Interestingly. at this radius also the transition between objects with relaxatior times smaller and larger than a Hubble time occurs (Fig. 13).," Interestingly, at this radius also the transition between objects with relaxation times smaller and larger than a Hubble time occurs (Fig. \ref{rad_ML}) )."
+ Globular clusters show a large scatter. and do clearly not aligi along the manifold.," Globular clusters show a large scatter, and do clearly not align along the manifold."
+ The third formulation includes all objects for fitting GCs and UCDs with log(r.)«-2.2., The third formulation includes all objects for fitting GCs and UCDs with $\rm log(r_e)<-2.2$.
+ This formulation hence excludes dynamically relaxed stellar systems from the fit., This formulation hence excludes dynamically relaxed stellar systems from the fit.
+ UCDs with log(r.)>—2.2 align very well with the manifold. while globular clusters and smaller UCDs do not.," UCDs with $\rm log(r_e)>-2.2$ align very well with the manifold, while globular clusters and smaller UCDs do not."
+ Summarizing. UCDs with log(r;)>-2.2 οz7 pe) appear to form a single family with larger stellar systems in the fundamental manifold.," Summarizing, UCDs with $\rm log(r_e)>-2.2$ $r_e\gtrsim$ 7 pc) appear to form a single family with larger stellar systems in the fundamental manifold."
+ The location of most GCs ts inconsistent with the fundamental manifold extrapolated from larger stellar systems., The location of most GCs is inconsistent with the fundamental manifold extrapolated from larger stellar systems.
+ Dynamically un-relaxed stellar systems appear to form a single manifold. while the relaxed systems — due to their advanced dynamical evolution — scatter very broadly around it.," Dynamically un-relaxed stellar systems appear to form a single manifold, while the relaxed systems — due to their advanced dynamical evolution — scatter very broadly around it."
+ In this paper we have analysed the internal dynamics of 23 ultra-compact dwarf galaxies in the Fornax cluster., In this paper we have analysed the internal dynamics of 23 ultra-compact dwarf galaxies in the Fornax cluster.
+ The analysis is based on high-resolution spectroscopy obtained with the FLAMES spectrograph at the VLT., The analysis is based on high-resolution spectroscopy obtained with the FLAMES spectrograph at the VLT.
+ Our targets cover an approximate mass range of 10°<M Me and a luminosity range -10.4<My-13.5 mag. overlapping the bright end of the globular cluster luminosity function.," Our targets cover an approximate mass range of $10^6<M<10^8$ $_{\sun}$ and a luminosity range $-10.4<M_V<-13.5$ mag, overlapping the bright end of the globular cluster luminosity function."
+ We also compare the dynamicalproperties of UCDs and GCs. and put them into the context of fundamental sealing relations defined for larger galaxies.," We also compare the dynamicalproperties of UCDs and GCs, and put them into the context of fundamental scaling relations defined for larger galaxies."
+ We obtain the following results:, We obtain the following results:
+the black hole masses in their sample of AGN observed with a slew of X-ray telescopes that include Ariel V. HEAO-1I. and the Einstein observatory.,"the black hole masses in their sample of AGN observed with a slew of X-ray telescopes that include Ariel V, HEAO-1, and the Einstein observatory."
+ Such a method however has been criticized for its dependence on data quality and coverage (Lawrence et al., Such a method however has been criticized for its dependence on data quality and coverage (Lawrence et al.
+ 1987: Mehardy Czerny 1987)., 1987; Mchardy Czerny 1987).
+ It does nevertheless provide a strong limit on the size of the emitting region., It does nevertheless provide a strong limit on the size of the emitting region.
+ For example. Reynolds et al. (," For example, Reynolds et al. ("
+1995) and Yaqoob et al. (,1995) and Yaqoob et al. (
+1997) have reported a factor of 1.5 increase in flux in as little as 100s in 6-30-15.,1997) have reported a factor of 1.5 increase in flux in as little as 100s in $-$ 6-30-15.
+ Others have resorted to employing power density spectrum (PDS) techniques (and variants thereof) as potential black hole mass estimators (e.g. Edelson Nandra 1999 for NGC3516: owak Chiang 1999 for 6-30-15: Hayashida et al., Others have resorted to employing power density spectrum (PDS) techniques (and variants thereof) as potential black hole mass estimators (e.g. Edelson Nandra 1999 for NGC3516; Nowak Chiang 1999 for $-$ 6-30-15; Hayashida et al.
+ 1998 introduces the for heir sample of AGN observed with Ginga)., 1998 introduces the for their sample of AGN observed with ).
+ However. even this is not without its caveats.," However, even this is not without its caveats."
+ X-ray variability studies of AGNs thus ur (with the exception of TIRAS18325-5926: Twasawa et al., X-ray variability studies of AGNs thus far (with the exception of IRAS18325-5926; Iwasawa et al.
+ 1998) jas shown that the power-law spectrum is essentially chaotic with no characteristic period., 1998) has shown that the power-law spectrum is essentially chaotic with no characteristic period.
+" The problem of unevenly sampled data streams. characteristic of X-ray observations of AGN. is overcome py the solution of Lomb (1976) and Scargle (1982: see Press et al 1992),"," The problem of unevenly sampled data streams, characteristic of X-ray observations of AGN, is overcome by the solution of Lomb (1976) and Scargle (1982; see Press et al 1992)."
+ Others still have searched for time lags and leads using cross correlation function techniques., Others still have searched for time lags and leads using cross correlation function techniques.
+ However. the problem associated with unevenly sampled data is still a concern. and has been addressed by e.g. Edelson Krolik (1988). Yaqoob et al (1997). Reynolds (1999). and in this paper.," However, the problem associated with unevenly sampled data is still a concern, and has been addressed by e.g. Edelson Krolik (1988), Yaqoob et al (1997), Reynolds (1999), and in this paper."
+ 6-30-15 is a bright nearby (220.0078) Seyfert | galaxy hat has been extensively studied by every major X-ray observatory since its identification., $-$ 6-30-15 is a bright nearby ) Seyfert 1 galaxy that has been extensively studied by every major X-ray observatory since its identification.
+ A recent study that takes advantage of the ugh energy and broad-band coverage of by Lee et al. (, A recent study that takes advantage of the high energy and broad-band coverage of by Lee et al. (
+1998. 1999) with respectively a SO ks observation in 1996. and 400 ks 1997 observation have confirmed the clear existence of a broad iron ine and reflection component in this object.,"1998, 1999) with respectively a 50 ks observation in 1996, and 400 ks 1997 observation have confirmed the clear existence of a broad iron line and reflection component in this object."
+ À previous study by Iwasawa et al. (, A previous study by Iwasawa et al. (
+1996) (hereafter I96) using a 4.5 day observation with revealed the iron line profile to be variable. which is confirmed in this long 1997| observation (Iwasawa et al.,"1996) (hereafter I96) using a 4.5 day observation with revealed the iron line profile to be variable, which is confirmed in this long 1997 observation (Iwasawa et al."
+" 1999, hereafter 199) which was observed simultaneously with."," 1999, hereafter I99) which was observed simultaneously with."
+ Additionally. Lee et al. (," Additionally, Lee et al. ("
+1999) have been able to constrain the iron abundance - reflection fraction relation for the first time using the observation.,1999) have been able to constrain the iron abundance - reflection fraction relation for the first time using the observation.
+ Guainazzi et al. (, Guainazzi et al. (
+1999) give gooc bounds on the high energy cutoff of the continuum from a observationof 6-30-15.,1999) give good bounds on the high energy cutoff of the continuum from a observation of $-$ 6-30-15.
+ We investigate in this paper changes in the direct and reflected components with our simultaneous and observations spanning a good time interval of zz 400 Ks. in order to speculate on ;xossible causes for variability.," We investigate in this paper changes in the direct and reflected components with our simultaneous and observations spanning a good time interval of $\approx$ 400 ks, in order to speculate on possible causes for variability."
+ In Section 3. we investigate rapic variability with colour ratios.," In Section 3, we investigate rapid variability with colour ratios."
+ This is followed in Section 4 by flux-correlated studies. and a detailed analysis of the time intervals surrounding the bright and flare shown in Fig. l..," This is followed in Section 4 by flux-correlated studies, and a detailed analysis of the time intervals surrounding the bright and flare shown in Fig. \ref{fig1-ascaxteltc}."
+ In Section 5. we shift gears to temporal studies and search (using cross-correlation techniques) for the presence of time lags and euds. in order to better understand the physical processes that may give rise to euch other.," In Section 5, we shift gears to temporal studies and search (using cross-correlation techniques) for the presence of time lags and leads, in order to better understand the physical processes that may give rise to each other."
+ We discuss in Section 6 tentative evidence for the presence of a possible 33 hour period. and break in he power spectrum.," We discuss in Section 6 tentative evidence for the presence of a possible 33 hour period, and break in the power spectrum."
+ Finally. we discuss the spectral and temporal tindings in Section 7.," Finally, we discuss the spectral and temporal findings in Section 7."
+ In particular. we discuss the challenge that spectral findings may pose to current models for reflection by cold material. and implications from temporal studies for placing constraints on the size of the emitting region and mass of the black hole in 6-30-15.," In particular, we discuss the challenge that spectral findings may pose to current models for reflection by cold material, and implications from temporal studies for placing constraints on the size of the emitting region and mass of the black hole in $-$ 6-30-15."
+ We present also a simple model to explain some of our enigmatic spectral findings., We present also a simple model to explain some of our enigmatic spectral findings.
+ We conclude in Section 8 with a summary of the pertinent findings from this study., We conclude in Section 8 with a summary of the pertinent findings from this study.
+ 6-30-15 was observed by over the period from 1997 Aug 4 to 1997 Aug I2 by both the Proportional Counter Array (PCA) and High-Energy X-ray Timing Experiment (HEXTE) instruments., $-$ 6-30-15 was observed by over the period from 1997 Aug 4 to 1997 Aug 12 by both the Proportional Counter Array (PCA) and High-Energy X-ray Timing Experiment (HEXTE) instruments.
+ We note that good data covered ~400 Ks. even though the observation spanned ~700 Ks.," We note that good data covered $\sim$ 400 ks, even though the observation spanned $\sim$ 700 ks."
+ Tt was simultaneously observed for ~200ks by the Solid-state Imaging Spectrometers (SIS) over the period 1997 August 3 to 1997 August 10 with a half-day gap part way through the observation., It was simultaneously observed for $\rm \sim 200~ks$ by the Solid-state Imaging Spectrometers (SIS) over the period 1997 August 3 to 1997 August 10 with a half-day gap part way through the observation.
+ We concentrate only on the PCA observations in this paper., We concentrate only on the PCA observations in this paper.
+ We extract PCA light curves and spectra from only the top Xenon layer using the software., We extract PCA light curves and spectra from only the top Xenon layer using the software.
+ The light curves were extracted with the default 16s time bins for Standard 2 PCA data time resolution., The light curves were extracted with the default 16s time bins for Standard 2 PCA data time resolution.
+ Data from PCUs 0. |. and 2 are combined to improve signal-to-noise at the expense of slightly blurring the spectral resolution.," Data from PCUs 0, 1, and 2 are combined to improve signal-to-noise at the expense of slightly blurring the spectral resolution."
+ Data from the remaining PCUs (PCU 3 and 4) are excluded because these instruments are known to periodically suffer discharge and are hence sometimes turned off., Data from the remaining PCUs (PCU 3 and 4) are excluded because these instruments are known to periodically suffer discharge and are hence sometimes turned off.
+ Good time intervals were selected to exclude any earth or South Atlantic Anomaly (SAA) passage occulations. and to ensure stable pointing.," Good time intervals were selected to exclude any earth or South Atlantic Anomaly (SAA) passage occulations, and to ensure stable pointing."
+ We also filter out electron contamination events., We also filter out electron contamination events.
+ We generate background data using in order to estimate the internal background caused by interactions between the radiation/particles and the detector/spacecraft at the time of observation., We generate background data using in order to estimate the internal background caused by interactions between the radiation/particles and the detector/spacecraft at the time of observation.
+ This is done by matching the conditions of observations with those in various model files., This is done by matching the conditions of observations with those in various model files.
+" The model files used are the L7-240 background models which are intended to be specialized for application to faint sources with count rate less than doctissμοι3,"," The model files used are the L7-240 background models which are intended to be specialized for application to faint sources with count rate less than 40 $\rm cts\,s^{-1}\,PCU^{-1}$."
+ The PCA response matrix for the data set was created usingV2., The PCA response matrix for the data set was created using.
+36. Background models and response matrices are representative of the most up-to-date PCA calibrations., Background models and response matrices are representative of the most up-to-date PCA calibrations.
+ Figure |. shows the background subtracted light curve for| and. with 256s binning.," Figure \ref{fig1-ascaxteltc} shows the background subtracted light curve for and, with 256s binning."
+ The time intervals corresponding to the and flares are decomposed further and analyzed in depth in Section 4.3., The time intervals corresponding to the and flares are decomposed further and analyzed in depth in Section 4.3.
+ Significant variability can be seen in both light curves on short and long timescales., Significant variability can be seen in both light curves on short and long timescales.
+ Flare and minima events are seen to correlate temporally in both light Curves., Flare and minima events are seen to correlate temporally in both light curves.
+ Intraday variability has been seen in most quasars and AGNs: in 6-30-15. rapid X-ray variability on the order of 100 s has been reported by a number of workers (e.g. Reynolds etal.," Intraday variability has been seen in most quasars and AGNs; in $-$ 6-30-15, rapid X-ray variability on the order of 100 s has been reported by a number of workers (e.g. Reynolds etal."
+ 1995: Yaqoob et al., 1995; Yaqoob et al.
+ 1997)., 1997).
+ These significant luminositychanges, These significant luminositychanges
+"We see that the tail and overrunning stream velocities are reasonably well reproduced with Mx 2. ος=1 and Mx|. eshc,=2.","We see that the tail and overrunning stream velocities are reasonably well reproduced with $M \approx 2$ , $c_{\rm S}/c_{\rm i} = 1$ and $M \approx 1$, $c_{\rm S}/c_{\rm i} = 2$."
+ Tails should show more complex velocity structure closer to the injection zone particularly with anisotropic mass injection but we do not consider this further here., Tails should show more complex velocity structure closer to the injection zone particularly with anisotropic mass injection but we do not consider this further here.
+ The setting up time for the tail of knot 38 =2-4xI0?yrs with a length of 2.5xI0em and a tail gas velocity of 10-20kms7!.," The setting up time for the tail of knot 38 $\approx 2-4 \times
+10^{3}\;{\rm yrs}$ with a length of $2.5 \times 10^{17}\;\cm$ and a tail gas velocity of $10-20 \;\kmps$."
+ This implies formation after fast wind activity has ceased., This implies formation after fast wind activity has ceased.
+ Although Patriarchi&Perinotto(1991). showed from IUE observations that of central stars of PNe emit fast particle winds. Cerruti-Sola&Perinotto(1985). failed to detect one from the central star of NGC 7293.," Although \citet{Patriarchi:1991} showed from IUE observations that of central stars of PNe emit fast particle winds, \citet{Cerruti-Sola:1985} failed to detect one from the central star of NGC 7293."
+" Shorter tails take proportionally shorter times to form and more ""typical"" tails with lengths of a few times 10em could be set up in timescales of the order of a hundred years.", Shorter tails take proportionally shorter times to form and more “typical” tails with lengths of a few times $10^{15}\;\cm$ could be set up in timescales of the order of a hundred years.
+ It is not clear why the tails are generally fairly short., It is not clear why the tails are generally fairly short.
+ One possibility might be time fluctuations in stream or source injection properties., One possibility might be time fluctuations in stream or source injection properties.
+ The effects on tail production mechanisms of systematic clump motions. radial or non-radial. have received little attention.," The effects on tail production mechanisms of systematic clump motions, radial or non-radial, have received little attention."
+ The non-radial component of the Helix knot expansion velocity is unclear., The non-radial component of the Helix knot expansion velocity is unclear.
+ On general grounds. non-radial velocity components of clumps seem inevitable if clumps form in dynamical instabilities.," On general grounds, non-radial velocity components of clumps seem inevitable if clumps form in dynamical instabilities."
+ If they are generated by shell fragmentation. they should have some dispersion around the shell velocity at fragmentation since it is highly unlikely that shells are strictly spherical.," If they are generated by shell fragmentation, they should have some dispersion around the shell velocity at fragmentation since it is highly unlikely that shells are strictly spherical."
+ Diamond&Kemball(2003) measured global proper motions of the SiO maser spots in the Mira variable TX Cam that correspond to velocities of about 10.kms! and noted that “significant” non-radial velocity components are present., \citet{Diamond:2003} measured global proper motions of the SiO maser spots in the Mira variable TX Cam that correspond to velocities of about $10\;\kmps$ and noted that “significant” non-radial velocity components are present.
+ These motions are in a stellar atmosphere of radius ~fewx1013 em., These motions are in a stellar atmosphere of radius $\sim {\rm few} \times 10^{13}\;\cm$ .
+ If these spots move out to radial distances ~10'em. conservation of angular momentum ensures very small radial velocity components (<|kms 5).," If these spots move out to radial distances $\sim 10^{17}\;\cm$, conservation of angular momentum ensures very small non-radial velocity components $\ll 1 \;\kmps$ )."
+ In Paper IV results are given for situations where clumps move at an angle to a global flow., In Paper IV results are given for situations where clumps move at an angle to a global flow.
+ To a good approximation. well-developed tails lie along the resultant velocity vector of the stream-clump velocities.," To a good approximation, well-developed tails lie along the resultant velocity vector of the stream-clump velocities."
+" If we assume that the stream Is strictly radial. then tails are aligned to an angle a=aretan[V44/(Vs—Ved] with the radial direction. where Vg, is the non-radial velocity component of the clump velocity."," If we assume that the stream is strictly radial, then tails are aligned to an angle $\alpha \approx {\rm
+arc\;\;tan}[V_{\rm tan}/(V_{\rm S}-V_{\rm C})]$ with the radial direction, where $V_{\rm tan}$ is the non-radial velocity component of the clump velocity."
+ In Table 3. we give maximum values of Vian that allow an alignment angle of 5° to the radial direction where we take this angle as characteristic.," In Table \ref{tab:vtan} we give maximum values of $V_{\rm
+tan}$ that allow an alignment angle of $^{\circ}$ to the radial direction where we take this angle as characteristic."
+" The best-fit cases for V4 and Vs detailed above require very small values of V4, (€2.kms7!) for alignment.", The best-fit cases for $V_{\rm T}$ and $V_{\rm S}$ detailed above require very small values of $V_{\rm tan}$ $\ltsimm 2\;\kmps$ ) for alignment.
+ This is unlikely to be satisfied with shell instability formation but is more consistent with the much lower non-radial velocity components expected if the clumps originated in stellar atmospheres., This is unlikely to be satisfied with shell instability formation but is more consistent with the much lower non-radial velocity components expected if the clumps originated in stellar atmospheres.
+ We have examined the persistence of tails where the injection Is switched off to quantify the proposal by Meaburnetal.(2005) that this is responsible for the production of headless spokes., We have examined the persistence of tails where the injection is switched off to quantify the proposal by \citet{Meaburn:2005} that this is responsible for the production of headless spokes.
+ In Fig., In Fig.
+ 5. we show time frames for the evolution of the tails produced for the two best fit cases in Section 4. (Le. Mx2. ος/α=] and MxlI. cs/e;=2).," \ref{fig:evolve} we show time frames for the evolution of the tails produced for the two best fit cases in Section \ref{sec:tailvz}
+ (i.e. $M \approx 2$, $c_{\rm S}/c_{\rm i}=1$ and $M \approx 1$, $c_{\rm S}/c_{\rm i}= 2$ )."
+ As before. for the case of a sound speed ratio of unity we give both the advected scalar plots and the density plots.," As before, for the case of a sound speed ratio of unity we give both the advected scalar plots and the density plots."
+ Tails persist over timescales of at least 50 or so (in dimensionless units)., Tails persist over timescales of at least 50 or so (in dimensionless units).
+ The actual time unit is equal to 2rc/Vi. where Vi is the gas velocity in the clump frame (1.9. Vs—Vc).," The actual time unit is equal to $2r_{\rm
+C}/V_{\rm F}$, where $V_{\rm F}$ is the gas velocity in the clump frame (i.e. $V_{\rm S}-V_{\rm C}$ )."
+ If we take the sound speed in the injected gas to be 11.7kms! as before. the flow speed in both cases is 23.4kms!.," If we take the sound speed in the injected gas to be $11.7\;\kmps$ as before, the flow speed in both cases is $23.4 \;\kmps$."
+ If we take the injection radius as about 10.em corresponding to a typical clump diameter. the tails persist for times at least ~1400.vr.," If we take the injection radius as about $10^{15}\;\cm$ corresponding to a typical clump diameter, the tails persist for times at least $\sim 1400\;{\rm yr}$."
+ It appears plausible therefore that headless tails are the remains of tails behind recently destroyed clumps., It appears plausible therefore that headless tails are the remains of tails behind recently destroyed clumps.
+ Linear structures have been observed in several planetary nebulae., Linear structures have been observed in several planetary nebulae.
+ This morphology has led to two classes of models., This morphology has led to two classes of models.
+ However morphology itself is not sufficient to discriminate between them., However morphology itself is not sufficient to discriminate between them.
+ Velocity structure is the key and only for two knots (14 and 38) in the Helix nebula ts the necessary data available., Velocity structure is the key and only for two knots (14 and 38) in the Helix nebula is the necessary data available.
+ The acceleration along the tail associated with the spectacular knot 38 suggests strongly that there is momentum transfer between a directed flow and a mass source., The acceleration along the tail associated with the spectacular knot 38 suggests strongly that there is momentum transfer between a directed flow and a mass source.
+ The advanced evolutionary state of the central star. and the details of the nebular structure. rule out any hypersonic stellar wind-mass source interaction.," The advanced evolutionary state of the central star, and the details of the nebular structure, rule out any hypersonic stellar wind-mass source interaction."
+ Although attractive in many respects. the interaction with a subsonic hot stream of shocked stellar wind gas is also ruled out because the existence of a reservoir of hot gas cannot be reconciled with the nebular structure.," Although attractive in many respects, the interaction with a subsonic hot stream of shocked stellar wind gas is also ruled out because the existence of a reservoir of hot gas cannot be reconciled with the nebular structure."
+ Much more plausible is tail formation taking place when a transonic or moderately supersonic stream of photoionized gas overruns gas injected by the photoionization of cometary globules., Much more plausible is tail formation taking place when a transonic or moderately supersonic stream of photoionized gas overruns gas injected by the photoionization of cometary globules.
+ The injected gas ts accelerated by momentum transfer from the stream and the velocities that are generated are in good agreement between those measured for the likely overrunning stream and the tail gas of knot 38 in the Helix., The injected gas is accelerated by momentum transfer from the stream and the velocities that are generated are in good agreement between those measured for the likely overrunning stream and the tail gas of knot 38 in the Helix.
+ Thus the general scheme for the formation of cometary globules involves flows (or winds) overrunninga slowly expanding system of pre- dense neutral globules., Thus the general scheme for the formation of cometary globules involves flows (or winds) overrunninga slowly expanding system of pre-existing dense neutral globules.
+ At least for NGC 7293. the knotty nature of the ambient neutral shells associated with this object surely precludes any other possibility (instabilities in a global expanding shell for example).," At least for NGC 7293, the knotty nature of the ambient neutral shells associated with this object surely precludes any other possibility (instabilities in a global expanding shell for example)."
+ In fact the very strong, In fact the very strong
+not cover many period eveles.,not cover many period cycles.
+ Figure 4. shows the Scaregle periodogram for the RN'PE/PCA data., Figure \ref{fig:756ls} shows the Lomb-Scargle periodogram for the RXTE/PCA data.
+ Independent confirmation that both the IN'TTIECIUCAL and the IUNTIS sources are probably SXP756 comes from 1e OGLE 1 optical monitoring of the counterpart., Independent confirmation that both the INTEGRAL and the RXTE sources are probably SXP756 comes from the OGLE III optical monitoring of the counterpart.
+ This source is unusual in that it shows striking optical behaviour wound the time of the periastron passage of the neutron gaar (Coe Edge. 2004).," This source is unusual in that it shows striking optical behaviour around the time of the periastron passage of the neutron star (Coe Edge, 2004)."
+ These authors quote an ephemeris for outbursts of T — JD2449830. | 394N. The outburst reported here peaked at JD2454945 which gives a phase for re INPEGRAL observations of 0.98 using the published ephemeris: coinciding with the sharp optical Hare seen in 1e OGLE 111 data at that time - see Figure 5.., These authors quote an ephemeris for outbursts of T = JD2449830 + 394N. The outburst reported here peaked at JD2454945 which gives a phase for the INTEGRAL observations of 0.98 using the published ephemeris; coinciding with the sharp optical flare seen in the OGLE III data at that time - see Figure \ref{fig:ogle1}.
+ Thus there can be no doubt that it is SXP756 that was seen in the INTEGRAL data., Thus there can be no doubt that it is SXP756 that was seen in the INTEGRAL data.
+ The significant neutron star spin up in this svstenm is ga10wn from our long-term IUNTIZ monitoring programme - see Figure 6.., The significant neutron star spin up in this system is shown from our long-term RXTE monitoring programme - see Figure \ref{fig:756period}.
+ Measurements up to MJD54000. come from Calache ct al 2008., Measurements up to MJD54000 come from Galache et al 2008.
+ From the cata shown in this figure it is possible to determine a straight line fit using à X72 routine andestablish an average 2 = -I128s/vear over —11 vears., From the data shown in this figure it is possible to determine a straight line fit using a $\chi^{2}$ routine andestablish an average $\dot{P}$ = -1.28s/year over $\sim$ 11 years.
+ Using ligure 5 of Coe. McBride Corbet (2010) this value implies a typical outhurst £2? of ~5 times that value.," Using Figure 5 of Coe, McBride Corbet (2010) this value implies a typical outburst $\dot{P}$ of $\sim$ 5 times that value."
+ The N-ray luminosity was calculated from the observed peak pulse amplitude in counts/pcu/s that occurred during the current outburst., The X-ray luminosity was calculated from the observed peak pulse amplitude in counts/pcu/s that occurred during the current outburst.
+ Vhis amplitude. is converted to luminosity assuming that 1 RXTE count/peu/s = 1077 erg/s at the distance of 60kpe to the SALC (though. the depth of the sources within the SAIC is unknown and allect this distance bv up to #£10kpe - Laney Stobie. 1986).," This amplitude is converted to luminosity assuming that 1 RXTE count/pcu/s = $\times 10^{37}$ erg/s at the distance of 60kpc to the SMC (though the depth of the sources within the SMC is unknown and affect this distance by up to $\pm$ 10kpc - Laney Stobie, 1986)."
+ Phe X-ray spectrum was assumed to be a power lay with a photon index = 1.5 and an Ny=6 πο..., The X-ray spectrum was assumed to be a power law with a photon index = 1.5 and an $N_{H}=6\times10^{20}cm^{-2}$ .
+ Furthermore it was assumect tha there was an average pulse, Furthermore it was assumed that there was an average pulse
+ , 
+accelerated in internal shocks (Figure 1)). and in the reverse shock curing transition to sell- expansion (Granot&Guetta2003).,"accelerated in internal shocks (Figure \ref{fig1}) ), and in the reverse shock during transition to self-similar expansion \citep{GnG}."
+.. As explained in the introduction. we consider the latter explanations less likely. since thev require modifications of the fireball physics (in particular. magnetic field well below ecquipartition in the fireball plasma). ancl since the temporal behavior is hard to account for in the internal shock scenario.," As explained in the introduction, we consider the latter explanations less likely, since they require modifications of the fireball physics (in particular, magnetic field well below equipartition in the fireball plasma), and since the temporal behavior is hard to account for in the internal shock scenario."
+ We have shown that the high energy data. when interpreted as emission from forward shock electrons. allow to put stringent constraints (Eq. 3))," We have shown that the high energy data, when interpreted as emission from forward shock electrons, allow to put stringent constraints (Eq. \ref{eq:constraints}) )"
+ on fireball parameters. such as the total (isotropic equivalent) enerev ZZ. initial Lorenz factor D; aud ambient medium density n.," on fireball parameters, such as the total (isotropic equivalent) energy $E$, initial Lorenz factor $\Gamma_i$ and ambient medium density $n$."
+ Such constraints can not be inferred from keVAleV observations alone., Such constraints can not be inferred from keV—MeV observations alone.
+ The inlerred values of E and D; are similar to those (vpical to cosmological GRBs. while the inferred constraint E/n>=10ουςcm implies a ratio E/n which is higher than tvpically obtained. Efín~10?lergcm? (Bloom.Frail&Kulkarni2003).," The inferred values of $E$ and $\Gamma_i$ are similar to those typical to cosmological GRBs, while the inferred constraint $E/n\gtrsim10^{56}{\rm erg\,cm}^3$ implies a ratio $E/n$ which is higher than typically obtained, $E/n\sim10^{54}{\rm erg\,cm}^3$ \citep{Bloom03}."
+" The inferred. values of €.p. €p,, ud €,., are consistent with (hose inferred from other GRD and afterglow observations."," The inferred values of $\epsilon_{e,f}$, $\epsilon_{B,r}$ and $\epsilon_{e,r}$ are consistent with those inferred from other GRB and afterglow observations."
+" The value of eg,p is usually less well constrained by observations. and is required to be well below equiparliGion in our case."," The value of $\epsilon_{B,f}$ is usually less well constrained by observations, and is required to be well below equipartition in our case."
+ This. ancl the large ratio of {η may account lor the rareness of GRBOLLOLT-tvpe high energy. (ails.," This, and the large ratio of $E/n$ may account for the rareness of GRB941017-type high energy tails."
+ The unknown recdshilt of the source of GRBOLLOLT is treated in our analysis as a [ree parameter: A choice of model parameters implies a choice of redshift (see ligure 1))., The unknown redshift of the source of GRB941017 is treated in our analysis as a free parameter: A choice of model parameters implies a choice of redshift (see figure \ref{fig1}) ).
+ Values of E close to the lower limit. c10°! erg. imply low redshift. 2~0.1.," Values of $E$ close to the lower limit, $\simeq10^{54}$ erg, imply low redshift, $z\sim0.1$."
+ This is a direct consequence “the fact that the gamma-ray [Inence of GIRDO4101T is unusually high. rather than of the fact that its high energy spectrum is unusually hard.," This is a direct consequence of the fact that the gamma-ray fluence of GRB941017 is unusually high, rather than of the fact that its high energy spectrum is unusually hard."
+ The fInence. 6.5x10.1'erg/cn. implies 1 (isotropic equivalent) gamma-ray energy release of ~107 erg [or z=1.," The fluence, $6.5\times10^{-4}{\rm erg/cm^2}$, implies an (isotropic equivalent) gamma-ray energy release of $\sim10^{55}$ erg for $z=1$."
+ Figure 1 demonstrates (hat dillerent scenarios accounting lor the hieh energy component X GRDB941017. as well as different model parameters in a given scenario. lead to different model predictions for the (Iuxes at optical. X-ray and sub-TeV energy bands.," Figure 1 demonstrates that different scenarios accounting for the high energy component of GRB941017, as well as different model parameters in a given scenario, lead to different model predictions for the fluxes at optical, X-ray and sub-TeV energy bands."
+ The predicted [Iuxes are well within the detection capabilities of the SWIFT (in optical) and GLAST (at 10100GeV) satellites. and of sub-TeV ground based Cerenkov telescopes (e.g. (MeEnery2003).. (Weekesefal...2002))). 3..," The predicted fluxes are well within the detection capabilities of the SWIFT (in optical) and GLAST (at 10–100GeV) satellites, and of sub-TeV ground based Cerenkov telescopes (e.g. \citep{McEnery03}, \citep{Weekes02}) \ref{eq:constraints}."
+Recent observations have shown the ubiquitous presence of propagating magnetohydrodynamie (MHD) Alfvénnic waves in the solar atmosphere.,Recent observations have shown the ubiquitous presence of propagating magnetohydrodynamic (MHD) Alfvénnic waves in the solar atmosphere.
+ For example. Alfvénnic transverse waves propagating in magnetic waveguides of the solar corona were first observed by Tomezyketal.(2007) and Tomezyk&MelIntosh(2009) using the Coronal Multichannel Polarimeter (CoMP). and more recently by McIntoshetal.(2011) using SDO/AIA.," For example, Alfvénnic transverse waves propagating in magnetic waveguides of the solar corona were first observed by \citet{tomczyk07} and \citet{tomczyk09} using the Coronal Multichannel Polarimeter (CoMP), and more recently by \citet{mcintosh2011} using SDO/AIA."
+ In chromospheric spicules. the presence of Doppler oscillations is known for more than 40 years (seethere-viewbyZaqarashvill&Erdélyi 2009).," In chromospheric spicules, the presence of Doppler oscillations is known for more than 40 years \citep[see the review by][]{temuryreview}."
+. Recent observations of Alfvénnic transverse waves in spicules have been reported by. e.g. DePontieuetal.(2007)... Zaqarashvilietal. (2007). Kimetal. (2008).. Heetal. (2009a.b).. Okamoto&DePon-tieu (2011).," Recent observations of Alfvénnic transverse waves in spicules have been reported by, e.g, \citet{depontieu07}, \citet{temuryspicules}, \citet{kim}, \citet{he1,he2}, \citet{okamotodepontieu}."
+. In addition. propagating transverse waves in thin threads of solar prominences have been observed (e.g..Linetal.2007.2009) and Alfvénnic waves in bright points have been reported (Jessetal.2009).," In addition, propagating transverse waves in thin threads of solar prominences have been observed \citep[e.g.,][]{lin07,lin09} and Alfvénnic waves in bright points have been reported \citep{jess2009}."
+. The role and implications of the observed Alfvennic waves for the heating of the solar atmosphere have been discussed by. e.g.. Erdélyi&FedunCargill&DeMoortel (2011).," The role and implications of the observed Alfvénnic waves for the heating of the solar atmosphere have been discussed by, e.g., \citet{robertus07,depontieu07,mcintosh2011,cargillineke}."
+ Based on MHD wave theory. a number of works have interpreted the observed waves as propagating kink MHD waves (e.g..Erdélyi&Fedun2007;VanDoorsselaereetal.al.2010;Verthet2010.2011:Soler201 1a.b.c)..," Based on MHD wave theory, a number of works have interpreted the observed waves as propagating kink MHD waves \citep[e.g.,][]{robertus07,tom08,lin09,pascoe1,pascoe2,TGV,VTG,verthspicule,solerspatial,resonantflow,stratified}."
+ Kink waves are transverse waves with mixed. fast MHD and Alfvénnic properties (see.e.g.Edwin&Roberts1983;Goossensetal. 2009).," Kink waves are transverse waves with mixed fast MHD and Alfvénnic properties \citep[see, e.g.,][]{edwinroberts,goossens09}."
+. In thin magnetic tubes kink waves are highly Alfvénnic because their dominant restoring force is magnetic tension (Goossensetal.2009)., In thin magnetic tubes kink waves are highly Alfvénnic because their dominant restoring force is magnetic tension \citep{goossens09}.
+. It has been shown that resonant absorption. caused by plasma inhomogeneity in the direction. transverse to the magnetic field. is a natural and efficient damping mechanism for kink waves (seetheGoossensetal. 2011).," It has been shown that resonant absorption, caused by plasma inhomogeneity in the direction transverse to the magnetic field, is a natural and efficient damping mechanism for kink waves \citep[see the recent reviews by][]{goossens06,goossens08,goossens11}."
+ In magnetic waveguides resonant absorption transfers wave energy from transverse kink motions to azimuthal motions localized in the inhomogeneous part of the waveguide., In magnetic waveguides resonant absorption transfers wave energy from transverse kink motions to azimuthal motions localized in the inhomogeneous part of the waveguide.
+ This process has been studied numerically by Pascoeetal.(2010.2011) in the case of driven kink waves in coronal waveguides.," This process has been studied numerically by \citet{pascoe1,pascoe2} in the case of driven kink waves in coronal waveguides."
+ Using analytical theory based on the thin tube and thin boundary approximations. Terradasetal.(2010.hereafterTGV) obtained that the damping length due to resonant absorption is inversely proportional to the wave frequency.," Using analytical theory based on the thin tube and thin boundary approximations, \citet[hereafter TGV]{TGV} obtained that the damping length due to resonant absorption is inversely proportional to the wave frequency."
+ Therefore. it was predicted that high-frequency kink waves become damped in length scales shorter than low-frequency waves.," Therefore, it was predicted that high-frequency kink waves become damped in length scales shorter than low-frequency waves."
+ Verthetal.(2010) showed that. this result Is consistentwith the CoMP observations of damped coronal waves (Tomezyketal.2007;Tomezyk&MelIntosh2009).," \citet{VTG} showed that this result is consistentwith the CoMP observations of damped coronal waves \citep{tomczyk07,tomczyk09}."
+. Subsequent investigations have extended the original work by TGV by incorporating effects not considered in their paper., Subsequent investigations have extended the original work by TGV by incorporating effects not considered in their paper.
+ For example. Soleretal.(2011b) took the presence of flow into account. and Soleretal.(2011e) studied the influence of longitudinal density stratification.," For example, \citet{resonantflow} took the presence of flow into account, and \citet{stratified} studied the influence of longitudinal density stratification."
+ Both works concluded that the damping length remains inversely proportional to the frequency when flows and longitudinal stratification are present., Both works concluded that the damping length remains inversely proportional to the frequency when flows and longitudinal stratification are present.
+ In TGV and subsequent works cited above. the plasma is assumed fully tonized.," In TGV and subsequent works cited above, the plasma is assumed fully ionized."
+ However. the plasma in the cooler parts of the solar atmosphere is only partially ionized as. e.g.. in the chromosphere or in prominences.," However, the plasma in the cooler parts of the solar atmosphere is only partially ionized as, e.g., in the chromosphere or in prominences."
+ This fact raises the relevant question on whether the damping length remains inversely proportional to the frequency when the plasma is partially ionized or. on the contrary. this dependence is modified by the effect of ton-neutral collisions.," This fact raises the relevant question on whether the damping length remains inversely proportional to the frequency when the plasma is partially ionized or, on the contrary, this dependence is modified by the effect of ion-neutral collisions."
+ The effect of partial ionization on the damping of Alfvénn waves has beer investigated m a large number of papers in different contexts., The effect of partial ionization on the damping of Alfvénn waves has been investigated in a large number of papers in different contexts.
+ Some examples are the works by. e.g.. Haerendel(1992):DePontieu&Haeren-(2010):Zaqarashvilietal. (2011a.b)..," Some examples are the works by, e.g., \citet{hearendel,depontieu98,pecseli,depontieu2001,forteza07,solerpartial,marc2010,temury,temuryhelium}."
+ However. these works only focus on the role of partial 10nizatior for the damping and do not consider resonant absorption. which is a basic and unavoidable phenomenon when plasma and/or magnetic inhomogeneity is present.," However, these works only focus on the role of partial ionization for the damping and do not consider resonant absorption, which is a basic and unavoidable phenomenon when plasma and/or magnetic inhomogeneity is present."
+ To our knowledge. the first attempt to study resonant waves in partially tonized plasmas was by Soleretal. (2009b)..," To our knowledge, the first attempt to study resonant waves in partially ionized plasmas was by \citet{solerpartialres}. ."
+ These authors used the single-fluid approximation (see.e.g..Bragin- to investigate standing resonant kink waves in à model of a partially ionized prominence thread.," These authors used the single-fluid approximation \citep[see, e.g.,][]{brag} to investigate standing resonant kink waves in a model of a partially ionized prominence thread."
+ Subsequently.," Subsequently,"
+profile is composed of a Gaussian and a Lorentzian of equal EWIIM. then a Gaussian fit overestimates the true EWIIM of the Voigt byLOY.,"profile is composed of a Gaussian and a Lorentzian of equal FWHM, then a Gaussian fit overestimates the true FWHM of the Voigt by."
+.. We find that the procedure of using Gaussian line widths in equation 2. to estimate the pressure broadening is a more reliable method than some alternatives., We find that the procedure of using Gaussian line widths in equation \ref{eq:voigtApprox} to estimate the pressure broadening is a more reliable method than some alternatives.
+ For example. one could imagine filling a Voigt profile civectly to an observed low frequency line {ο determine simultaneously both the Gaussian and Lorentzian components.," For example, one could imagine fitting a Voigt profile directly to an observed low frequency line to determine simultaneously both the Gaussian and Lorentzian components."
+ Llowever. we find (his unreliable because the fitting procedure assigns the line width to either the Gaussian or Lorentzian components primarilv on the shape of the line wines where the profiles differ most strongly. but the signal-to-noise ratio is (he weakest.," However, we find this unreliable because the fitting procedure assigns the line width to either the Gaussian or Lorentzian components primarily on the shape of the line wings where the profiles differ most strongly, but the signal-to-noise ratio is the weakest."
+ Alternatively. one might estimate the FWIIM directly from the 3 channels with the peak and hall-power emission.," Alternatively, one might estimate the FWHM directly from the 3 channels with the peak and half-power emission."
+ However. because this estimate relies on single channel measurements it is less reliable than a procedure that uses measurements at all the channels across (he line prolile.," However, because this estimate relies on single channel measurements it is less reliable than a line-fitting procedure that uses measurements at all the channels across the line profile."
+ several lines of evidence. discussed later in (he paper suggest (hat there are density variations within the III regions.," Several lines of evidence, discussed later in the paper suggest that there are density variations within the HII regions."
+ Because all the III regions in our sample are unresolved. (he measured electron density represents the average within the II] region.," Because all the HII regions in our sample are unresolved, the measured electron density represents the average within the HII region."
+ The size of the IU] region within the beam (filling factor) does not affect this estimate of the average density because the electron densities are determined [rom the line widths rather (han the emission intensity., The size of the HII region within the beam (filling factor) does not affect this estimate of the average density because the electron densities are determined from the line widths rather than the emission intensity.
+ However. the unknown density structure of the ILLI region creates some uncertainty.," However, the unknown density structure of the HII region creates some uncertainty."
+ If there are significant variations in optical depth with frequency then it is possible that lines al different frequencies could be generated in different locations within the IJ] reeion., If there are significant variations in optical depth with frequency then it is possible that lines at different frequencies could be generated in different locations within the HII region.
+ Aside from the svstematic effects. the uncertainty in the average electron densitv is given bv the standard propagation of errors. assuming the errors in the widths are independent.," Aside from the systematic effects, the uncertainty in the average electron density is given by the standard propagation of errors, assuming the errors in the widths are independent."
+ Uncertainties in the electron densities are listed in table 5. and lor the line center velocities and widths in table 4.., Uncertainties in the electron densities are listed in table \ref{electrondensities} and for the line center velocities and widths in table \ref{CMrrl}.
+ At the high densities in our HII regions. the centimeter continuum emission is optically thick to [ree-[ree emission.," At the high densities in our HII regions, the centimeter continuum emission is optically thick to free-free emission."
+" This can be determined from the formula llenderson1967) or (equ.10.Keto2003).. where T; is the electron temperature and 71, is the electron density squared times the path length."," This can be determined from the formula \citep[eqn. A.1b,][]{MezgerHenderson1967}
+ or \citep[eqn. 10,][]{Keto2003}, where $T_e$ is the electron temperature and $n_e^2L$ is the electron density squared times the path length."
+ For a characteristic size of 0.01 pe. an HII region has an optical depth of unity," For a characteristic size of 0.01 pc, an HII region has an optical depth of unity"
+Horizontal-branch stars are evolved intermediate-mass stars that are burning helium in their core (e.g. Hoyle Schwarzschild 1955).,Horizontal-branch stars are evolved intermediate-mass stars that are burning helium in their core (e.g. Hoyle Schwarzschild 1955).
+ Some hot blue horizontal-branch (hereafter BHB) stars with Ty larger than approximately 11500 K are very interesting objects since they exhibit several observational anomalies., Some hot blue horizontal-branch (hereafter BHB) stars with $T_{\rm eff}$ larger than approximately 11500 K are very interesting objects since they exhibit several observational anomalies.
+ This paper aims to study certain aspects of these astronomical objects in the light of recent modelling and observational results., This paper aims to study certain aspects of these astronomical objects in the light of recent modelling and observational results.
+ First. abundance anomalies are observed for such BHB stars in several globular clusters (Glaspey et al.," First, abundance anomalies are observed for such BHB stars in several globular clusters (Glaspey et al."
+ 1989: Behr et al., 1989; Behr et al.
+ 1999: Moehler et al., 1999; Moehler et al.
+ 1999: Behr et al., 1999; Behr et al.
+ 20002: Behr 2003a: Hubrig et al., 2000a; Behr 2003a; Hubrig et al.
+ 2009)., 2009).
+ Khalack et al. (, Khalack et al. (
+2007. 2008 and 2010) have detected vertical stratification of the abundance of several chemical elements including Fe in the atmosphere of some stars of this type.,"2007, 2008 and 2010) have detected vertical stratification of the abundance of several chemical elements including Fe in the atmosphere of some stars of this type."
+ Khalack et al. (, Khalack et al. (
+2008) found that the iron abundance increases toward the lower atmosphere in three BHB stars.,2008) found that the iron abundance increases toward the lower atmosphere in three BHB stars.
+ Nitrogen and. sulfur stratification was also detected in the hot BHB star HD 135485 (Khalack et al., Nitrogen and sulfur stratification was also detected in the hot BHB star HD 135485 (Khalack et al.
+ 2007)., 2007).
+ The abundances of these two elements are found to increase toward the upper atmosphere., The abundances of these two elements are found to increase toward the upper atmosphere.
+ More recentlv. Khalack et al. (," More recently, Khalack et al. ("
+2010) determined the vertical Fe abundance gradient in a sample of 14 BHB stars using the observed spectra. of Behr (2003:).,2010) determined the vertical Fe abundance gradient in a sample of 14 BHB stars using the observed spectra of Behr (2003a).
+ They detected vertical Fe stratification in five (and possibly seven) of these BHB stars., They detected vertical Fe stratification in five (and possibly seven) of these BHB stars.
+ Isotopic anomalies have also been recently detected in BHB stars (Hubrig et al., Isotopic anomalies have also been recently detected in BHB stars (Hubrig et al.
+ 2009)., 2009).
+ Photometric jumps are observed on the hot side of Ziv 11500 K for the horizontal-branch sequence of several globular clusters (Grundahl et al., Photometric jumps are observed on the hot side of $T_{\rm eff} \simeq $ 11500 K for the horizontal-branch sequence of several globular clusters (Grundahl et al.
+ 1999) as compared to what is predicted by canonical models., 1999) as compared to what is predicted by canonical models.
+ Photometric gaps are also observed for this sequence (Ferraro et al., Photometric gaps are also observed for this sequence (Ferraro et al.
+ 1998) at the same Ίωῃ where the photometric jump Mentioned above occurs., 1998) at the same $T_{\rm eff}$ where the photometric jump mentioned above occurs.
+ Another anomaly detected is that the rotational velocities of BHB stars are observed to drop abruptly for 72; approximately larger than 11500 K (Peterson. Rood Crocker 1995: Behr et al.," Another anomaly detected is that the rotational velocities of BHB stars are observed to drop abruptly for $T_{\rm eff}$ approximately larger than 11500 K (Peterson, Rood Crocker 1995; Behr et al."
+ 2000a and 2000b: Behr 2003b)., 2000a and 2000b; Behr 2003b).
+ Finally. spectroscopic gravities of BHB stars in. globular clusters. at least for the metal-poor ones (e.g. Crocker et al.," Finally, spectroscopic gravities of BHB stars in globular clusters, at least for the metal-poor ones (e.g. Crocker et al."
+ 1988: Moehler et al., 1988; Moehler et al.
+ 1995). are also lower than those predicted by classical BHB models.," 1995), are also lower than those predicted by classical BHB models."
+ Atomic diffusion (Michaud 1970) is effective only if the medium is hydrodynamically stable enough to prevent mixing. because of the order of magnitude of the diffusion velocities which are much smaller than those of macroscopic motions.," Atomic diffusion (Michaud 1970) is effective only if the medium is hydrodynamically stable enough to prevent mixing, because of the order of magnitude of the diffusion velocities which are much smaller than those of macroscopic motions."
+ In the case of BHB stars. their relatively low rotational velocities should render their superficial layer stable enough for diffusion to take place.," In the case of BHB stars, their relatively low rotational velocities should render their superficial layer stable enough for diffusion to take place."
+ For instance. recently. Quievy et al. (," For instance, recently, Quievy et al. ("
+2009) showed that for values of rotational velocities found for BHB stars with Z;y above 11500 K. meridional circulation in not efficient enough to prevent He,"2009) showed that for values of rotational velocities found for BHB stars with $T_{\rm eff}$ above 11500 K, meridional circulation in not efficient enough to prevent He"
+12 ! feature has disappeared. and that feature A is a new feature which has flared up not [ar from the 12 ! position.,"12 $^{-1}$ feature has disappeared, and that feature A is a new feature which has flared up not far from the 12 $^{-1}$ position."
+ On the other hand the identification of feature A with the 12 ! feature is supported bv its dominance in both spectra and by the equality of flux densities., On the other hand the identification of feature A with the 12 $^{-1}$ feature is supported by its dominance in both spectra and by the equality of flux densities.
+ All maser spots have been partially resolved., All maser spots have been partially resolved.
+ From Table 1 one can see that the largest spols are spot Ix. (20 masx 12 mas) and spot N (14 masx 3 mas). but they may consist of several smaller spots which we can not resolve.," From Table 1 one can see that the largest spots are spot K (20 $\times$ 12 mas) and spot N (14 $\times$ 3 mas), but they may consist of several smaller spots which we can not resolve."
+ The rest of the spots can be fitted with Gaussians wilh the major axis from 4.4 (o 10.7 mas. and the minor axis from 0.7 to 4.6 mas.," The rest of the spots can be fitted with Gaussians with the major axis from 4.4 to 10.7 mas, and the minor axis from 0.7 to 4.6 mas."
+ Considering that the beam size is 12 x 4 mas. 0.7 mas must be regarded only as a formal fitting solution.," Considering that the beam size is 12 $\times$ 4 mas, 0.7 mas must be regarded only as a formal fitting solution."
+ The real extent of the minor axis can be much less than the indicated. so it should be consklered as an upper limit.," The real extent of the minor axis can be much less than the indicated, so it should be considered as an upper limit."
+ The brightness temperature of one of the brightest 1665 MlIIz spots. F (reference feature). with the integrated flux density of 18.1 Jv is 1.6x 107 IX. The brightest 1667 MIIz spot is L. with the integrated Πιν density 31.9 Jv. and it has a brightness temperature of 1.9x 107 IX. These must be regarded as lower limits since most of the spots are probably not resolved along the minor axis.," The brightness temperature of one of the brightest 1665 MHz spots, F (reference feature), with the integrated flux density of 18.1 Jy is $1.6\times$ $^{12}$ K. The brightest 1667 MHz spot is L, with the integrated flux density 31.9 Jy, and it has a brightness temperature of $\times$ $^{12}$ K. These must be regarded as lower limits since most of the spots are probably not resolved along the minor axis."
+ All the mapped spots are elongated. wilh the major to minor axis ratio 3 or higher.," All the mapped spots are elongated, with the major to minor axis ratio 3 or higher."
+ If (he size of the maser spots is not intrinsic and is caused by the scattering in the interstellar medium. then the elongation of the spots may be a result of the anisotropic scattering.," If the size of the maser spots is not intrinsic and is caused by the scattering in the interstellar medium, then the elongation of the spots may be a result of the anisotropic scattering."
+ All maser spots given in Table 1 were imaged in all Stokes parameters., All maser spots given in Table 1 were imaged in all Stokes parameters.
+ This made ib possible to determine lull polarization properties of the maser spots., This made it possible to determine full polarization properties of the maser spots.
+ The percentage of circular polarization given in Table 1 was calculated as Positive V corresponds to the right circular polarization., The percentage of circular polarization given in Table 1 was calculated as Positive $V$ corresponds to the right circular polarization.
+ Percentage of the linear polarization was caleulated as: and the position angle of the electric vector of linear polarization was calculated as, Percentage of the linear polarization was calculated as: and the position angle of the electric vector of linear polarization was calculated as
+"lx magnitudes in Figure 9.. with completeness limits and lines ofconstant 6,dy overlaid.","K magnitudes in Figure \ref{bjvsK}, with completeness limits and lines ofconstant $b_J - K$ overlaid."
+ The plot in Figure 9 revealed where in 5; and Ix. magnitude parameter space the number of quasar detections declines: and. by extension revealed where the 5b;/—AN distribution will drop oll due to selection effects.," The plot in Figure \ref{bjvsK} revealed where in $b_J$ and K magnitude parameter space the number of quasar detections declines; and, by extension revealed where the $b_J - K$ distribution will drop off due to selection effects."
+ The upper left region of Figure 9. is where we expect the least quasar detections due to the combination of completeness. and the magnitude limits of the 56; and Ix xhotometry. referred to collectively as selection. elfects.," The upper left region of Figure \ref{bjvsK} is where we expect the least quasar detections due to the combination of completeness, and the magnitude limits of the $b_J$ and K photometry, referred to collectively as selection effects."
+ A ow quasars are found with 6;—dy4 and only one is ound satisbving A«17 and 5;20.5., A few quasars are found with $b_J - K > 4$ and only one is found satisfying $K < 17$ and $b_J > 20.5$.
+ As the decrease in observed. quasar detections coincides with where it isexpected to decrease because of selection elfects: we conclude rom Figure 9 that the decrease in quasar detections. redder han b;dv=4. is à result of sample selection and not an intrinsic property of our quasar sample.," As the decrease in observed quasar detections coincides with where it isexpected to decrease because of selection effects; we conclude from Figure \ref{bjvsK} that the decrease in quasar detections, redder than $ b_J - K = 4$, is a result of sample selection and not an intrinsic property of our quasar sample."
+ This is consistent with the fact that redcder quasars than those found here have en observed. such as in Clikmanetal.(2004)..," This is consistent with the fact that redder quasars than those found here have been observed, such as in \citet{2004ApJ...607...60G}."
+" The analysis of our sample selection supports the ivpothesis that the 6,dy distribution. extends. further and rededer than observed: here. as such our observed. red quasar fraction is a robust lower limit of the actual red quasar fraction for Ax;18.4."," The analysis of our sample selection supports the hypothesis that the $b_J - K$ distribution extends further and redder than observed here, as such our observed red quasar fraction is a robust lower limit of the actual red quasar fraction for $K \leq 18.4$."
+ Weighting the quasars by the completeness of the FCSS. using the weights in Table L.. we calculated the observed. red quasar fraction and found that a robust lower limit for the actual red. quasar fraction is ~0.22.," Weighting the quasars by the completeness of the FCSS, using the weights in Table \ref{weightingtable}, we calculated the observed red quasar fraction and found that a robust lower limit for the actual red quasar fraction is $\sim$ 0.22."
+ ‘To estimate the red. quasar fraction for Ax;18.4. we first assumed. that the b;9A colour distribution is independent of IX magnitude.," To estimate the red quasar fraction for $K \leq 18.4$, we first assumed that the $b_J - K$ colour distribution is independent of K magnitude."
+" Our second. assumption in estimating the red quasar fraction. was that the true extent of the bydv distribution of radio-quiet. quasars should match that of radio-loud quasars: therefore. we assumed that the truce extent of the b;—A distribution for our sample is by,dy Ξ δ. matching the extent of the 6,Av distribution of the PLLJES radio-oud quasars in Websteretal.(1995)."," Our second assumption in estimating the red quasar fraction, was that the true extent of the $b_J - K$ distribution of radio-quiet quasars should match that of radio-loud quasars; therefore, we assumed that the true extent of the $b_J - K$ distribution for our sample is $b_J - K$ = 8, matching the extent of the $b_J - K$ distribution of the PHJFS radio-loud quasars in \citet{1995Natur.375..469W}."
+. Using these two assumptions. we postulated a model for the number density of quasars in b;—A vs. Ix parameter space. then used. maximum likelihoodestimation to find values for the model parameters: and finally integrated this mocel to estimate the red. quasar fraction.," Using these two assumptions, we postulated a model for the number density of quasars in $b_J - K$ vs. K parameter space, then used maximum likelihoodestimation to find values for the model parameters; and finally integrated this model to estimate the red quasar fraction."
+ As Figure 4. and the results in IHichardsetal.(2003). and Llopkinsetal.(2004) are consistent with a b;A colour distribution. where red quasars occupy a tail extending— away from. blue. quasars. [or the by—.A. component of our model we used a Gamma distribution.," As Figure \ref{bjMKhistogram} and the results in \citet{2003AJ....126.1131R} and \citet{2004AJ....128.1112H} are consistent with a $b_J - K$ colour distribution, where red quasars occupy a tail extending away from blue quasars, for the $b_J - K$ component of our model we used a Gamma distribution."
+ For the IX. component of our model we have assumed a power-law. 1057.," For the K component of our model we have assumed a power-law, $10^{K\beta}$."
+ Normalised by the number of quasars in our sample. the resultant model for the number density of quasars in;ἐν vs. IX parameter space is. and maximum likelihood estimates of the mioclel parameters are ϐ=46.65. 39=0.39. a=0.36 and 5=8.56.," Normalised by the number of quasars in our sample, the resultant model for the number density of quasars in $b_J - K$ vs. K parameter space is, and maximum likelihood estimates of the model parameters are $\theta = 46.65$, $\beta = 0.39$, $\alpha = 0.36$ and $\gamma = 8.56$."
+ Due to limited cata. we were unable to test the validity of our moclel in describing the data statisiticallv: therefore. we assessed the model validity by plotting the observed b;fy colour and Ix magnitude distribution with that predicted by our model in Figure 10.. and comparing them.," Due to limited data, we were unable to test the validity of our model in describing the data statisitically; therefore, we assessed the model validity by plotting the observed $b_J - K$ colour and K magnitude distribution with that predicted by our model in Figure \ref{ModelTest}, , and comparing them."
+ Comparing the observed. ancl predicted. distributions in Figure 10.. we concluded that the model was suitable for estimating the red quasar fraction: because it reproduced the overall shape and scale of both the Ix and 5;ἐνdistributions.," Comparing the observed and predicted distributions in Figure \ref{ModelTest}, , we concluded that the model was suitable for estimating the red quasar fraction; because it reproduced the overall shape and scale of both the K and $b_J - K$distributions."
+" Integrating our model over 15xly 18.4. for the 6,A ranges Osb;|Nx Sand 3.5xb;Lv S. we calculated"," Integrating our model over $15 \leq K \leq 18.4$ , for the $b_J - K$ ranges $0 \leq b_J - K \leq 8$ and $3.5 \leq b_J - K \leq 8$ , we calculated"
+"the XMMI data with an absorbed power-law, but require an additional soft component for the XMM2 data.","the XMM1 data with an absorbed power-law, but require an additional soft component for the XMM2 data."
+" Our fits suggest that the XMM1 1984)spectrum is more typical of a steep power-law state, while the XMM2 spectrum is consistent with a thermally dominated state."," Our fits suggest that the XMM1 spectrum is more typical of a steep power-law state, while the XMM2 spectrum is consistent with a thermally dominated state."
+" A large [κο in the fit with SIMPL is problematic for our accretion disk model, which does not account for the effects of irradiation by neighboring corona."," A large $f_{\rm sc}$ in the fit with SIMPL is problematic for our accretion disk model, which does not account for the effects of irradiation by neighboring corona."
+" Hence, we only report our results from the XMM2 analysis below."," Hence, we only report our results from the XMM2 analysis below."
+ We note that the XMM1 data generally favor a best-fit M that is higher than XMM2 for the same i., We note that the XMM1 data generally favor a best-fit $M$ that is higher than XMM2 for the same $i$.
+" This yields a cooler disk model, consistent with the fact that more of the high energy flux is accounted for with the SIMPL component in these data."," This yields a cooler disk model, consistent with the fact that more of the high energy flux is accounted for with the SIMPL component in these data."
+" However, the joint confidence contours on a, and M are much larger for the XMM1 data, so the best-fit values are still consistent with the XMM?2 values at confidence."," However, the joint confidence contours on $a_*$ and $M$ are much larger for the XMM1 data, so the best-fit values are still consistent with the XMM2 values at confidence."
+" 'TheSwift andChandra observations caught the source at higher luminosities than the XMM2 observation, but the overall signal-to-noise is still lower because of the lower effective area of theSwift XRT and the short duration (10 ks) of theChandra ACIS-S observation."," The and observations caught the source at higher luminosities than the XMM2 observation, but the overall signal-to-noise is still lower because of the lower effective area of the XRT and the short duration (10 ks) of the ACIS-S observation."
+" Due to the lower signal-to-noise and stronger disk dominance, a suitable fit is provided by the BHSPEC model alone."," Due to the lower signal-to-noise and stronger disk dominance, a suitable fit is provided by the (absorbed) BHSPEC model alone."
+ The addition of the SIMPL (absorbed)model only provides a marginal improvement to the fit and the best-fit SIMPL parameters are poorly constrained., The addition of the SIMPL model only provides a marginal improvement to the fit and the best-fit SIMPL parameters are poorly constrained.
+ We also fit the combined S1 and $2 observations., We also fit the combined S1 and S2 observations.
+" In this case we tie all BHSPEC parameters together, except for 6, which is allowed to vary independently for S1 and S2."," In this case we tie all BHSPEC parameters together, except for $\ell$, which is allowed to vary independently for S1 and S2."
+" Even for the combined dataset, we find a good fit with BHSPEC alone, and poor constraints on the SIMPL model parameters."," Even for the combined dataset, we find a good fit with BHSPEC alone, and poor constraints on the SIMPL model parameters."
+" Therefore, we do not include the SIMPL model in subsequent analysis of the combined S1 and S2 datasets or in our analysis of theChandra data."," Therefore, we do not include the SIMPL model in subsequent analysis of the combined S1 and S2 datasets or in our analysis of the data."
+ We first consider the XMM2 observation., We first consider the XMM2 observation.
+ The soft thermal component is largely characterized by only two parameters: the energy at which the spectrum peaks and the overall normalization., The soft thermal component is largely characterized by only two parameters: the energy at which the spectrum peaks and the overall normalization.
+" In practice, this leads to degeneracies in the best-fit parameters of BHSPEC2006),, unless all but two of the parameters(see can be independently constrained."," In practice, this leads to degeneracies in the best-fit parameters of BHSPEC, unless all but two of the parameters can be independently constrained."
+" Since we can only constrain D independently, we will need to consider joint variation of the remaining parameters."," Since we can only constrain $D$ independently, we will need to consider joint variation of the remaining parameters."
+" For completeness, our best-fit parameters for various choices of i or a, are summarized in Table 1.."," For completeness, our best-fit parameters for various choices of $i$ or $a_*$ are summarized in Table \ref{t:xmm}."
+ We find acceptable fits for all i., We find acceptable fits for all $i$.
+ The best-fit model and data for i—0 are plotted in Figure 1.., The best-fit model and data for $i=0$ are plotted in Figure \ref{f:spec}.
+" We report the confidence interval for a single parameter, but due to the significant degeneracies in the model, the joint confidence contours better illustrate the actual parameter uncertainties."," We report the confidence interval for a single parameter, but due to the significant degeneracies in the model, the joint confidence contours better illustrate the actual parameter uncertainties."
+ Hence they are the focus of this work., Hence they are the focus of this work.
+" Since M and a, are of primary interest to us, we first examine their joint confidence contours, and consider the joint confidence of M and i in Section 4.."," Since $M$ and $a_*$ are of primary interest to us, we first examine their joint confidence contours, and consider the joint confidence of $M$ and $i$ in Section \ref{disc}."
+" For illustration, it is useful to consider several sets of contours for different choices of (fixed) 1, leaving £ as a free parameter."," For illustration, it is useful to consider several sets of contours for different choices of (fixed) $i$ , leaving $\ell$ as a free parameter."
+ These best-fit joint confidence contours are shown in Figure 2.., These best-fit joint confidence contours are shown in Figure \ref{f:spin}.
+" We consider five choices for cos£, evenly spaced from 0 to 1."," We consider five choices for $\cos i$, evenly spaced from 0 to 1."
+" Each set of contours has three curves corresponding to6896,,9096,, and confidence."," Each set of contours has three curves corresponding to, and confidence."
+ These are determined by the change in x? relative to the best-fit values listed in Table 1.., These are determined by the change in $\chi^2$ relative to the best-fit values listed in Table \ref{t:xmm}.
+" Even for fixed i, there is a clear correlation of a, and M."," Even for fixed $i$, there is a clear correlation of $a_*$ and $M$."
+" At confidence, the entire range of a, (—1«à,« 0.99) is allowed,and the corresponding M varies by a factor of 4-5 over this range."," At confidence, the entire range of $a_*$ $-1 < a_* < 0.99$ ) is allowed,and the corresponding $M$ varies by a factor of 4-5 over this range."
+" The only exception is for nearly edge on systems (cosi~ 0), for which low spins are disallowed."," The only exception is for nearly edge on systems $\cos i \sim
+0$ ), for which low spins are disallowed."
+" The correlation of best-fit M with 7 is very strong as well, consistentwith previous modeling of other ULX sources 2008)."," The correlation of best-fit $M$ with $i$ is very strong as well, consistentwith previous modeling of other ULX sources ."
+". This is illustrated more clearly in Figure 3,, which shows"," This is illustrated more clearly in Figure \ref{f:inc}, , which shows"
+only four are at 2κ 2: 222?) while the remaining DLAs without detectable 21 em absorption all have 30 lower limits of >700 K on the spin temperature.,"only four are at $z \gtrsim 2$ ; \citealt{wolfe79,wolfe81,kanekar06,kanekar07}) ) while the remaining DLAs without detectable 21 cm absorption all have $3\sigma$ lower limits of $> 700$ K on the spin temperature."
+ The preponderance of high spin temperature estimates in 2 DLAs has usually been attributed to a relatively low fraction of the cold neutral phase of (the CNM). with most of the gas in the warm phase (the WNM) (e.g. 229).," The preponderance of high spin temperature estimates in $z$ DLAs has usually been attributed to a relatively low fraction of the cold neutral phase of (the CNM), with most of the gas in the warm phase (the WNM) (e.g. \citealt{carilli96,chengalur00}) )."
+ This assumes that the absorbers have a similar two-phase structure to that seen in the ISM of the Milky Way (e.g. 25., This assumes that the absorbers have a similar two-phase structure to that seen in the ISM of the Milky Way (e.g. \citealt{wolfe03b}) ).
+ The low spin temperature of the 2~2.289 DLA towards TXS 03112430 would then imply that a higher fraction of the aalong the line of sight is in the CNM compared with the majority of DLAs at =2., The low spin temperature of the $z \sim 2.289$ DLA towards TXS 0311+430 would then imply that a higher fraction of the along the line of sight is in the CNM compared with the majority of DLAs at $z \gtrsim 2$.
+ ? argued that the high derived 7; values of high-z DLAs could be due to the low metallicities of typical > DLAs. with the paucity of metals resulting in fewer radiation pathways for gas cooling (See also 2)).," \citet{chengalur00}
+ argued that the high derived $\ts$ values of $z$ DLAs could be due to the low metallicities of typical $z$ DLAs, with the paucity of metals resulting in fewer radiation pathways for gas cooling (see also \citealt{young97}) )."
+ If this is correct. one woulc expect DLAs with low spin temperatures (such as the absorber towards TXS 03114430) to have significantly higher metallicities than those of the general DLA population (2)..," If this is correct, one would expect DLAs with low spin temperatures (such as the absorber towards TXS 0311+430) to have significantly higher metallicities than those of the general DLA population \citep{kanekar01a}."
+ As expected. and assuming that the ALSOS equivalent width has not been over-estimated. the 52.289 DLA has [Si/H] 0.45. one of mos metal-rich DLAs yet discovered.," As expected, and assuming that the $\lambda 1808$ equivalent width has not been over-estimated, the $z \sim 2.289$ DLA has [Si/H] $\ge -0.48$, one of most metal-rich DLAs yet discovered."
+ It has also been suggested that the observed low 21 em optica depth in high-z DLAs arises due to covering factor effects (e.g. 23)., It has also been suggested that the observed low 21 cm optical depth in $z$ DLAs arises due to covering factor effects (e.g. \citealt{curran05}) ).
+ In the present case. any reduction in the covering factor below the assumed value of unity can only strengthen the case for a low spin temperature (since 7.=(138+36)> f£.Ky.," In the present case, any reduction in the covering factor below the assumed value of unity can only strengthen the case for a low spin temperature (since $\ts = (138 \pm 36) \times
+f$ K)."
+ The only way to alter the conclusion that a sizeable fraction of aalong the line of sight is in the cold phase is if there are large spatial differences in the ccolumn densities along the optical and radio lines of sight (e.g. ?))., The only way to alter the conclusion that a sizeable fraction of along the line of sight is in the cold phase is if there are large spatial differences in the column densities along the optical and radio lines of sight (e.g. \citealt{wolfe03b}) ).
+ For example. if the optical QSO lies behind a “hole” in the ccolumn density distribution. the average ccolumn density against the radio QSO could be larger than that measured from the Lyman-a line. implying a higher spin temperature from Eqn. |..," For example, if the optical QSO lies behind a “hole” in the column density distribution, the average column density against the radio QSO could be larger than that measured from the $\alpha$ line, implying a higher spin temperature from Eqn. \ref{eqn:tspin}."
+ Unfortunately. it is very difficult to directly test this possibility.," Unfortunately, it is very difficult to directly test this possibility."
+ While Galactic ccolumn densities derived from Lyman-a absorption studies are in excellent agreement with those obtained from 22] em emission observations in the same directions. despite the very different spatial resolutions in the two methods. such comparisons have only been carried out for a fairly small number of high latitude lines of sight (2)..," While Galactic column densities derived from $\alpha$ absorption studies are in excellent agreement with those obtained from 21 cm emission observations in the same directions, despite the very different spatial resolutions in the two methods, such comparisons have only been carried out for a fairly small number of high latitude lines of sight \citep{dickey90}."
+ A similar comparison between vvalues derived from Lyman-a absorption and 21 em emission has only been possible in one DLA. the 2~0.009 absorber towards SBS 15434593.," A similar comparison between values derived from $\alpha$ absorption and 21 cm emission has only been possible in one DLA, the $z \sim 0.009$ absorber towards SBS 1543+593."
+ Here. the ccolumn densities agree to within a factor of ~2. despite the extremely poor spatial resolution (5.3h.l Kpe) of the radio observations (2).," Here, the column densities agree to within a factor of $\sim 2$, despite the extremely poor spatial resolution $\sim 5.3 \: h_{71}^{-1}$ kpc) of the radio observations \citep{chengalur02}."
+. While both of these studies suggest that the vvalues along the radio and optical lines of sight are likely to be comparable. we cannot formally rule out the possibility of differences in an individual absorber.," While both of these studies suggest that the values along the radio and optical lines of sight are likely to be comparable, we cannot formally rule out the possibility of differences in an individual absorber."
+ However. the fact that the expected high metallicity is indeed seen in the +~2.289 DLA (assuming that our Si AISO08 measurement is accurate) is consistent with the interpretation of a high CNM fraction.," However, the fact that the expected high metallicity is indeed seen in the $z \sim 2.289$ DLA (assuming that our Si $\lambda$ 1808 measurement is accurate) is consistent with the interpretation of a high CNM fraction."
+ Finally. ? noted a relationship between spin temperature and absorber morphology. in that. at low redshifts. low 7. values are only found in DLAs identified with luminous disk galaxies. while low-z. low-luminosity DLAs. associated with dwarf or LSB galaxies. are all found to have high spin temperatures K).," Finally, \citet{kanekar03} noted a relationship between spin temperature and absorber morphology, in that, at low redshifts, low $\ts$ values are only found in DLAs identified with luminous disk galaxies, while $z$, low-luminosity DLAs, associated with dwarf or LSB galaxies, are all found to have high spin temperatures K)."
+ It has not been possible to test this empirical relationship üt z>1. ας the host galaxies of high-z DLAs are rarely detectable.," It has not been possible to test this empirical relationship at $z > 1$, as the host galaxies of $z$ DLAs are rarely detectable."
+ However. if the relationship does extend out to high redshifts. we would expect the 2~2.289 DLA to be a massive. luminous disk galaxy.," However, if the relationship does extend out to high redshifts, we would expect the $z \sim 2.289$ DLA to be a massive, luminous disk galaxy."
+ We note that both the high metallicity and the large velocity spreads seen in the optical low-ionization metal lines and the 21em absorption are consistent with the absorption arising in a massive galaxy., We note that both the high metallicity and the large velocity spreads seen in the optical low-ionization metal lines and the 21cm absorption are consistent with the absorption arising in a massive galaxy.
+ If so. it should be possible to detect the absorber host with deep imaging: this would be the first direct test of the T.-morphology relationship at high redshifts.," If so, it should be possible to detect the absorber host with deep imaging; this would be the first direct test of the $\ts$ -morphology relationship at high redshifts."
+ In summary. we have detected damped Lyman-a and 21 em absorption at 2.=2.289 towards the quasar TXS 03114430.," In summary, we have detected damped $\alpha$ and 21 cm absorption at $z = 2.289$ towards the quasar TXS 0311+430."
+ We obtain a DLA spin temperature of 7;= K. the first case of a low spin temperature estimate in a high redshift DLA.," We obtain a DLA spin temperature of $\ts = (138 \pm 36) \times f$ K, the first case of a low spin temperature estimate in a high redshift DLA."
+ The low spin temperature. high metallicity and large velocity spread of the 21 em and metallines all suggest that the absorber is likely to be a massive disk galaxy.," The low spin temperature, high metallicity and large velocity spread of the 21 cm and metallines all suggest that the absorber is likely to be a massive disk galaxy."
+Since the early surveys with the Uhuru and OSO-7 satellites it has been known that X-ray sources occur about à thousand times more frequently in globular clusters (GCs) than in the rest of the Galaxy (e.g..Katz1975:Clark1975).,"Since the early surveys with the Uhuru and OSO-7 satellites it has been known that X-ray sources occur about a thousand times more frequently in globular clusters (GCs) than in the rest of the Galaxy \citep[e.g.,][]{jk75,gc75}."
+. Over fifteen hundred of them have been detected to date in these stellar systems by recent missions such as Chandra and XMM-Newton (Pooley2010.andreferencestherein)..., Over fifteen hundred of them have been detected to date in these stellar systems by recent missions such as Chandra and XMM-Newton \citep[][ and references therein]{poo10}.
+ In addition. over a hundred millisecond pulsars — objects which have evolved in low-mass X-ray binaries — are known to reside in GCs (Lynchetal.2010).," In addition, over a hundred millisecond pulsars – objects which have evolved in low-mass X-ray binaries – are known to reside in GCs \citep{lyn10}."
+" Most of the X-ray sources are binaries with degenerate components (neutron stars or white dwarfs). henceforth referred to as ""degenerate binaries”."," Most of the X-ray sources are binaries with degenerate components (neutron stars or white dwarfs), henceforth referred to as “degenerate binaries”."
+ Remarkably. no stellar-mass black holes (BH) are known in Galactic GCs. despite theoretical predictions that they should exist (e.g...Deveechietal.2007).," Remarkably, no stellar-mass black holes (BH) are known in Galactic GCs, despite theoretical predictions that they should exist \citep[e.g., ][]{dev07}."
+. The presence of intermediate-mass black holes (IMBHs) (100M.1077 M.) in GCs has also not yet been proven (e.g..Mac-carone&Servillat2008;vanderMarelAnderson 2010).," The presence of intermediate-mass black holes (IMBHs) $100\,M_{\odot}< M <10^{3-4}~M_{\odot}$ ) in GCs has also not yet been proven \citep[e.g.,][]{mac08,van10}."
+. While the origin of IMBHs ts unclear. stellar-mass black holes should form along with neutron stars via supernova explosions during the early evolution of a cluster.," While the origin of IMBHs is unclear, stellar-mass black holes should form along with neutron stars via supernova explosions during the early evolution of a cluster."
+ The very presence of such objects would interesting information concerning the effectiveness of dynamical interactions leading to the ejection of BHs into the intracluster medium. which according to the predictions of current approximate models should be very high (e.g..Downingetal.2010).," The very presence of such objects would interesting information concerning the effectiveness of dynamical interactions leading to the ejection of BHs into the intracluster medium, which according to the predictions of current approximate models should be very high \citep[e.g.,][]{dow10}."
+. Their absence would confirm those predictions., Their absence would confirm those predictions.
+ Alternatively. it could mean that massive low-metallicity stars cannot produce BHs. which would have important implications for the origin of early BHs. believed to be seeds for the first galaxies in the Universe.," Alternatively, it could mean that massive low-metallicity stars cannot produce BHs, which would have important implications for the origin of early BHs, believed to be seeds for the first galaxies in the Universe."
+ In either case. searching for black holes in GCs is certainly a worthwhile task.," In either case, searching for black holes in GCs is certainly a worthwhile task."
+ It is well known that field X-ray novae (binaries most probably hosting stellar-mass black holes) spend most of their time in quiescence. showing only ellipsoidal variability in the visible domain (Remillard&MeClintock2006).," It is well known that field X-ray novae (binaries most probably hosting stellar-mass black holes) spend most of their time in quiescence, showing only ellipsoidal variability in the visible domain \citep{rem06}."
+.. Since GCs harbor a multitude of degenerate binaries. they must also contain many systems which presently accrete at a very low rate or do not acerete at all. thus being very weak or undetectable in the X-rays.," Since GCs harbor a multitude of degenerate binaries, they must also contain many systems which presently accrete at a very low rate or do not accrete at all, thus being very weak or undetectable in the X-rays."
+ Based on X-ray luminosity (Lx ) and hardness ratio. Grindlay (2006) splits the population of weak X-ray sources in GCs into four major classes.," Based on X-ray luminosity $L_{\mathrm X}$ ) and hardness ratio, Grindlay (2006) splits the population of weak X-ray sources in GCs into four major classes."
+ The first three. arranged according to Ly increasing from 107? (the present sensitivity limit for the nearest GCs) to 107 erg s! are i) active binaries (binary main-sequence stars in which the X-ray emission originates from chromospheric activity. e.g. BY Dra type systems): i) cataclysmic variables (in which a white dwarf accretes from a low-mass main-sequence companion); in) quiescent low-mass X-ray binaries (in which a neutron star intermittently accretes from a main-sequence or evolved companion). hereafter referred to as qLMXBs.," The first three, arranged according to $L_X$ increasing from $10^{29}$ (the present sensitivity limit for the nearest GCs) to $10^{32}$ erg $^{-1}$ are i) active binaries (binary main-sequence stars in which the X-ray emission originates from chromospheric activity, e.g. BY Dra type systems); ii) cataclysmic variables (in which a white dwarf accretes from a low-mass main-sequence companion); iii) quiescent low-mass X-ray binaries (in which a neutron star intermittently accretes from a main-sequence or evolved companion), hereafter referred to as qLMXBs."
+ The fourth class is composed of millisecond pulsars. whose luminosity strongly depends on the predominant mechanism generating X-ray quanta (residual accretion: collision of the pulsar wind with the ambient medium or with matter lost by the companion: thermal surface emission).," The fourth class is composed of millisecond pulsars, whose luminosity strongly depends on the predominant mechanism generating X-ray quanta (residual accretion; collision of the pulsar wind with the ambient medium or with matter lost by the companion; thermal surface emission)."
+ The study of weak X-ray sources 15 hampered by the fact that even in the nearest GCs they are detected down to the present sensitivity limit. and often remain unclassified as they are too dim to allow for an estimate of the hardness ratio.," The study of weak X-ray sources is hampered by the fact that even in the nearest GCs they are detected down to the present sensitivity limit, and often remain unclassified as they are too dim to allow for an estimate of the hardness ratio."
+ Most of them are likely active binaries or cataclysmic variables. but some may be quiescent systems with neutron stars or ever black holes.," Most of them are likely active binaries or cataclysmic variables, but some may be quiescent systems with neutron stars or even black holes."
+ Optical counterparts of these systems are ofter weak and hard to identify in the crowded GC environment (e.g..Verbuntetal. 2008).," Optical counterparts of these systems are often weak and hard to identify in the crowded GC environment \citep[e.g.,][]{ver08}."
+. So far. the X-ray data have allowed us to identify a few candidate GLMXBs. but their nature has not beer confirmed by optical spectroscopy (Guillotetal.2009).," So far, the X-ray data have allowed us to identify a few candidate qLMXBs, but their nature has not been confirmed by optical spectroscopy \citep{gui09}."
+. We decided to follow an entirely different approach., We decided to follow an entirely different approach.
+ Instead of looking for optical counterparts of X-ray sources. we select a sample of candidate systems for degenerate binaries from short-period. low-amplitude optical variables cataloguec in existing surveys. and measure their radial velocities.," Instead of looking for optical counterparts of X-ray sources, we select a sample of candidate systems for degenerate binaries from short-period, low-amplitude optical variables catalogued in existing surveys, and measure their radial velocities."
+ Our targets have nearly sinusoidal light curves with V-banc amplitudes smaller than 0.35 mag and periods shorter thar ~1.3 days., Our targets have nearly sinusoidal light curves with $V$ -band amplitudes smaller than 0.35 mag and periods shorter than $\sim$ 1.3 days.
+ This type of variability is common in all classes of close degenerate binaries: in particular. it is observed i1," This type of variability is common in all classes of close degenerate binaries; in particular, it is observed in"
+their model showed that the emission wings originate in regions that are much closer to the star than the radiation that is forming the upper part of the line.,their model showed that the emission wings originate in regions that are much closer to the star than the radiation that is forming the upper part of the line.
+ Moreover. we also attempted to measure the RVs of other available spectral linesA.. and A)) to see if they undergo similar time changes.," Moreover, we also attempted to measure the RVs of other available spectral lines, and ) to see if they undergo similar time changes."
+ We primarily measured also the outer emission wings of these lines. but for the line it was possible to obtain a relatively accurate RV of the central absorption core.," We primarily measured also the outer emission wings of these lines, but for the line it was possible to obtain a relatively accurate RV of the central absorption core."
+ The RVs measured on the emission wings of the and lines were averaged., The RVs measured on the emission wings of the and lines were averaged.
+ Two methods of RV measurement were used as follows., Two methods of RV measurement were used as follows.
+ We denote the RVs measured by the first method asmanual and those measured by the second method asautomatic to distinguish them in following sections., We denote the RVs measured by the first method as and those measured by the second method as to distinguish them in following sections.
+ In Fig., In Fig.
+ 5 the automatic emission RVs are plotted vs. the manually measured ones to see whether there Is any systematic difference between the two methods., \ref{compall} the automatic emission RVs are plotted vs. the manually measured ones to see whether there is any systematic difference between the two methods.
+ We fitted the data with a linear relation and somewhat surprisingly the slope was found to be 0.968+40.009: i.e. theautomaticmethod Findsaslightlynarrowertotalrange em," We fitted the data with a linear relation and somewhat surprisingly the slope was found to be $0.968$ $0.009$; i.e., the automatic method finds a slightly narrower total range of RV variations than the manual one."
+ission are published in detail in Table 2 for the manual. and in Table 3 for the automatic The emission RVs measured manually and automatically are plotted vs. time in Fig. 6...," All individual RV measurements on the steep wings of the emission are published in detail in Table 2 for the manual, and in Table 3 for the automatic The emission RVs measured manually and automatically are plotted vs. time in Fig. \ref{13hahe}."
+ Additional time plots for other measured features can be found in Appendix AppendixA:.., Additional time plots for other measured features can be found in Appendix \ref{apen}.
+ Figure 6 shows that exhibits long-term RV variations over several years. which seem to correlate with those of the peak intensity of the emission.," Figure \ref{13hahe} shows that exhibits long-term RV variations over several years, which seem to correlate with those of the peak intensity of the emission."
+ To be able to search for periodic RV changes on a shorter timescale. one has first to remove the long-term ones.," To be able to search for periodic RV changes on a shorter timescale, one has first to remove the long-term ones."
+ To check how robust the result is or how much it depends on the specifiο way of secular-changes removal. we applied three differer= approaches to this goal.," To check how robust the result is or how much it depends on the specific way of secular-changes removal, we applied three different approaches to this goal."
+ The first was to use the program HECI3 written. by PH. which is based on the smoothing technique developed by ?? and which uses some subroutines kindly provided by Dr.Vondrákl?.. The level of smoothing is controlled by a smoothing parameter e (the lower the value of e. the higher the smoothing). and the smoothing routine can operate either through individual data points or through suitably chosen normal points. which are the weighted mean values of the observed quantity (RV in our case) over the chosen constant time intervals.," The first was to use the program HEC13 written by PH, which is based on the smoothing technique developed by \citet{vondrak1969, vondrak1977} and which uses some subroutines kindly provided by Dr. The level of smoothing is controlled by a smoothing parameter $\epsilon$ (the lower the value of $\epsilon$, the higher the smoothing), and the smoothing routine can operate either through individual data points or through suitably chosen normal points, which are the weighted mean values of the observed quantity (RV in our case) over the chosen constant time intervals."
+ In both cases. the rresidua are provided for all individual observations.," In both cases, the residua are provided for all individual observations."
+ In these, In these
+It is now widely accepted that the large-scale structure we see in the distribution of galaxies arises through a process of gravitational instability. which amplifies primordial fluctuations laid down in the very early Universe.,"It is now widely accepted that the large-scale structure we see in the distribution of galaxies arises through a process of gravitational instability, which amplifies primordial fluctuations laid down in the very early Universe."
+ The rate at which structure grows from these small perturbations offers a key discriminant between cosmological models., The rate at which structure grows from these small perturbations offers a key discriminant between cosmological models.
+ For instance. dark energy models in which general relativity is unmoditied predict different Large-Scale Structure formation compared with Moditied Gravity models with the same background expansion (e.g.Dvalietal.2000:CarrollBrans2005:Nesseris&Perivolaropoulos 2008).," For instance, dark energy models in which general relativity is unmodified predict different Large-Scale Structure formation compared with Modified Gravity models with the same background expansion \citep[e.g.][]{dvali00,carroll04,Bra05,NesPer08}."
+. Structure growth is driven by the motion of matter (and inhibited by the cosmological expansion)., Structure growth is driven by the motion of matter (and inhibited by the cosmological expansion).
+ Galaxies are expected to act as test particles within this matter flow. so the motion of galaxies carries an imprint of the rate of growth of large-scale structure.," Galaxies are expected to act as test particles within this matter flow, so the motion of galaxies carries an imprint of the rate of growth of large-scale structure."
+ Because of this. many previous analyses have shown that observations of these galaxy peculiar velocities can distinguish between classes of models (e.g.Jain&Zhang2007:Song&Koyama2008:Percival 2008)..," Because of this, many previous analyses have shown that observations of these galaxy peculiar velocities can distinguish between classes of models \citep[e.g.][]{jain07,song08,song08b}."
+. A key technique to statistically measure the growth of the velocity field. uses redshift-space distortions seen in galaxy surveys (Kaiser1987)., A key technique to statistically measure the growth of the velocity field uses redshift-space distortions seen in galaxy surveys \citep{Kai87}.
+. Galaxy maps. produced by estimating distances from redshifts obtained in spectroscopic galaxy surveys. reveal an anisotropic galaxy distribution.," Galaxy maps, produced by estimating distances from redshifts obtained in spectroscopic galaxy surveys, reveal an anisotropic galaxy distribution."
+ The anisotropies arise because galaxy recession velocities. from which distances are inferred. include components from. both the Hubble flow and peculiar velocities from the comoving motions of galaxies.," The anisotropies arise because galaxy recession velocities, from which distances are inferred, include components from both the Hubble flow and peculiar velocities from the comoving motions of galaxies."
+ These distortions. encode information about the build-up of structure., These distortions encode information about the build-up of structure.
+ Many previous surveys have been analysed to measure O'/b. where h is the deterministic. local. linear bias of the galaxies.," Many previous surveys have been analysed to measure $\beta\approx\Omega_m^{0.6}/b$ , where $b$ is the deterministic, local, linear bias of the galaxies."
+ The latest generation of large surveys have providec ever tighter constraints., The latest generation of large surveys have provided ever tighter constraints.
+ Analyses using the 2-degree Field Galaxy Redshift Survey (2dFGRS:Collessetal.2003). have measurec redshift-space distortions in both the correlation function etal.2001:Hawkins2003). and power spectrum etal. 2004).," Analyses using the 2-degree Field Galaxy Redshift Survey \citep[2dFGRS;][]{colless03} have measured redshift-space distortions in both the correlation function \citep{peacock01,hawkins03} and power spectrum \citep{percival04}."
+. Using the Sloan Digital Sky Survey (SDSS:Yorketal. 2000). redshift-space distortions have also been measured in the correlation function (Zehavietal.2005:Okumuraeal.2008:Cabré&Gaztanaga 2008). and using an Eigenmode decomposition to separate real and redshift-space effects (Tegmarketal.2004. 2006)..," Using the Sloan Digital Sky Survey \citep[SDSS;][]{york00}, redshift-space distortions have also been measured in the correlation function \citep{Zeh05,Oku08,cabre08}, , and using an Eigenmode decomposition to separate real and redshift-space effects \citep{tegmark04,tegmark06}. ."
+ Thesestudies were recently extended to z5| (Guzzoetal.2007) using the VIMOS-VLT Deep Survey (VVDS:LeFevreetal.2005:Garilli2008) and the 2SLAQ survey (daAngelaetal.2008).," Thesestudies were recently extended to $z\simeq1$ \citep{guzzo08} using the VIMOS-VLT Deep Survey \citep[VVDS;][]{lefevre05,garilli08} and the 2SLAQ survey \citep{Ang08}."
+. In addition to measuring B at z= 0.8. Le emphasised theimportance of using large-scalepeculiar velocitiesfor constraining models of cosmic acceleration.," In addition to measuring $\beta$ at $z=0.8$ , \citet{lefevre05} emphasised theimportance of using large-scalepeculiar velocitiesfor constraining models of cosmic acceleration."
+ On linear scales the theory behind the observed redshift-space distortions is well developed (Kaiser1987:Hamilton 1998)... and," On linear scales the theory behind the observed redshift-space distortions is well developed \citep{Kai87,HamiltonReview}, , and"
+Weak gravitational lensing is one of the most promising observational tools available to cosmologists for studying the recent accelerated expansion of the Universe.,Weak gravitational lensing is one of the most promising observational tools available to cosmologists for studying the recent accelerated expansion of the Universe.
+ Images of distant galaxies appear distorted due to the bending of light as it passes through the gravitational potential of the intervening matter., Images of distant galaxies appear distorted due to the bending of light as it passes through the gravitational potential of the intervening matter.
+ The image distortion can be described as a shearing of the original galaxy shape., The image distortion can be described as a shearing of the original galaxy shape.
+ This “cosmic shear” can be exploited to study the matter distribution in the Universe and the growth of structure., This “cosmic shear” can be exploited to study the matter distribution in the Universe and the growth of structure.
+ Cosmic shear has the potential to be one of the most powerful probes both of dark matter and of dark energy which together are thought to make up about 95% of the energy budget of the Universe (e.g.etal.," Cosmic shear has the potential to be one of the most powerful probes both of dark matter and of dark energy \citep{esoesa,detf} which together are thought to make up about $95\%$ of the energy budget of the Universe \citep[e.g.][]{spergelea06}."
+ For any [2007)..individual galaxy it is impossible to separate the small cosmic shear distortion from the intrinsic ellipticity of the galaxy shape., For any individual galaxy it is impossible to separate the small cosmic shear distortion from the intrinsic ellipticity of the galaxy shape.
+" However, the light from physically close galaxies will follow a similar trajectory to the observer, passing through spacetime curved by the same gravitational fields."," However, the light from physically close galaxies will follow a similar trajectory to the observer, passing through spacetime curved by the same gravitational fields."
+ Therefore these galaxies will acquire the same cosmic shear distortion., Therefore these galaxies will acquire the same cosmic shear distortion.
+ If it is assumed that galaxy shape and orientation are randomly assigned across the sky then the cosmic shear signal can be retrieved by averaging the ellipticities of a number of galaxies close on the sky., If it is assumed that galaxy shape and orientation are randomly assigned across the sky then the cosmic shear signal can be retrieved by averaging the ellipticities of a number of galaxies close on the sky.
+ In practice this assumption of randomly distributed galaxy shapes is unrealistic., In practice this assumption of randomly distributed galaxy shapes is unrealistic.
+" Spatially localised galaxies are expected to have formed within the same large-scale gravitational field, which is likely to cause an alignment in their intrinsic ellipticities (Heavens&Peacock||1988; Aubertal,))2004)."," Spatially localised galaxies are expected to have formed within the same large-scale gravitational field, which is likely to cause an alignment in their intrinsic ellipticities \citep{heavensp88,catelankb01,jing2002,aubert_pichon_colombi_2004}."
+. See [Schafer(2008) for a recent review., See \cite{schaefer_2008_review} for a recent review.
+" This etintrinsic alignment (IA) effect will appear as a systematic error in estimates of cosmological parameters extracted from cosmic shear data, unless accurately taken into account."," This intrinsic alignment (IA) effect will appear as a systematic error in estimates of cosmological parameters extracted from cosmic shear data, unless accurately taken into account."
+ Two types of IA affect the measured cosmic shear signal., Two types of IA affect the measured cosmic shear signal.
+ Physically close galaxies form in the same large- gravitational potential so share a preferred ellipticity, Physically close galaxies form in the same large-scale gravitational potential so share a preferred ellipticity
+"quark matter changes significantly,",quark matter changes significantly.
+ The strange stars become more compact. the magnetic field reducing the mass and radius of the star.," The strange stars become more compact, the magnetic field reducing the mass and radius of the star."
+ The surface energy. and the curvature term of the quark phase both diverge. with the surface tension diverging logarithmically. while the curvature term diverges much faster.," The surface energy and the curvature term of the quark phase both diverge, with the surface tension diverging logarithmically, while the curvature term diverges much faster."
+ Therefore the thermal nucleation of quark bubbles in a compact metastable state of neutron matter is completely forbidden in the presence of a strong magnetic field., Therefore the thermal nucleation of quark bubbles in a compact metastable state of neutron matter is completely forbidden in the presence of a strong magnetic field.
+ These results for the formation of quark bubbles in neutron matter have been obtained bv using the Csernad and Ixapusta (theory of nucleation and ils extension including the thermal conductivitv of dense matter 1994)., These results for the formation of quark bubbles in neutron matter have been obtained by using the Csernai and Kapusta theory of nucleation and its extension including the thermal conductivity of dense matter .
+. In this theory both the hadron aud quark materials are considered as substances with zero barvonic number. ancl are treated by using a relativistic formalism.," In this theory both the hadron and quark materials are considered as substances with zero baryonic number, and are treated by using a relativistic formalism."
+ The principal result in (his approach is the suggestion that the prefactor 7 1s proportional to the transport coellicients (viscosity and thermal conductivitv) of the neutron matter. and thus when these coellicients vanish a quark bubble necessarily does not form.," The principal result in this approach is the suggestion that the prefactor $k$ is proportional to the transport coefficients (viscosity and thermal conductivity) of the neutron matter, and thus when these coefficients vanish a quark bubble necessarily does not form."
+ llowever. argued that the energy flow does not vanish in absence of any heat conduction or viscous damping.," However, argued that the energy flow does not vanish in absence of any heat conduction or viscous damping."
+ Since the change of energy density € in time is given. in the low velocity limit. bv the conservation equation Je/0l=—V(er) 1992a).. where ie is (he enthalpy and ο the velocity. (his implies that the enerey flow ~ic is always present.," Since the change of energy density $e$ in time is given, in the low velocity limit, by the conservation equation $\partial e/\partial t=-\nabla \left( w\vec{v}%
+\right) , where $w$ is the enthalpy and $\vec{v}$ the velocity, this implies that the energy flow $\sim w\vec{v}$ is always present."
+ Therefore. an expression for the prefactor can be derived. which does not vanish in the absence of viscosity.," Therefore, an expression for the prefactor can be derived, which does not vanish in the absence of viscosity."
+ The viscous effects cause only a small perturbation to the prefactor., The viscous effects cause only a small perturbation to the prefactor.
+ The differences between the Csernai-Ixapusta (CHI) and Buggieri-Friedinan (RF) results are due to the technical differences in (he treatment of the pressure gradients., The differences between the Csernai-Kapusta (CK) and Ruggieri-Friedman (RF) results are due to the technical differences in the treatment of the pressure gradients.
+ A generalized approach. following the Cly formalism leading to the prefactor in both viscous and non-viscous regimes was developed by(2001).," A generalized approach, following the CK formalism leading to the prefactor in both viscous and non-viscous regimes was developed by."
+. Unlike in(1992a).. the linearized relativistic hydrodynamics equations have been solved in all regions.," Unlike in, the linearized relativistic hydrodynamics equations have been solved in all regions."
+ It is the purpose of the present paper to extend the Zeldovich nucleation theory. for the formation of quark droplets in neutron matter. by taking into account both the effects of the enerey flow and of thermal conductivity. shear and bulk viscosities.," It is the purpose of the present paper to extend the Zel'dovich nucleation theory for the formation of quark droplets in neutron matter, by taking into account both the effects of the energy flow and of thermal conductivity, shear and bulk viscosities."
+ As a result a more general expression for the quark matter droplet rate lormation can be obtained. which is also valid in the limiting case of vanishing conductivity and. viscosity coefficients.," As a result a more general expression for the quark matter droplet rate formation can be obtained, which is also valid in the limiting case of vanishing conductivity and viscosity coefficients."
+ In the nucleation theory of(1992a).. the (ransiGion from the neutron to quark state is possible only for a normal matter having viscous properties.," In the nucleation theory of, the transition from the neutron to quark state is possible only for a normal matter having viscous properties."
+ For zero bulk ancl shear viscosilies the transition is impossible., For zero bulk and shear viscosities the transition is impossible.
+ This assumption is. however. to restrictive. and the transition can also take place for a perfect (no viscosity or heat conduction)neutron matter," This assumption is, however, to restrictive, and the transition can also take place for a perfect (no viscosity or heat conduction)neutron matter"
+hese simulations to test ti6 WSB shear measurement methocl.,these simulations to test the KSB shear measurement method.
+ Overall. we find that his method is rather accurate. out with several provisos: we find a residual anti-correlation rctween the PSE ellipticity τind the corrected. ellipticities of faint galaxies.," Overall, we find that this method is rather accurate, but with several provisos: we find a residual anti-correlation between the PSF ellipticity and the corrected ellipticities of faint galaxies."
+ This effect. can be mace negligible if zünt galaxies (with SYNἘν 15) are removed from the catalogue., This effect can be made negligible if faint galaxies (with $S/N \la 15$ ) are removed from the catalogue.
+ We also find that the recovered. shear is linearly related to the input shear. but with a coellicient. of abou LS which must be used to calibrate the final shear.," We also find that the recovered shear is linearly related to the input shear, but with a coefficient of about 0.8 which must be used to calibrate the final shear."
+ With these precautions. the KSB method is sullicient for he current weak lensing surveys.," With these precautions, the KSB method is sufficient for the current weak lensing surveys."
+ However. the moethoc is neither optimal nor necessarily extendable to superior observing conditions.," However, the method is neither optimal nor necessarily extendable to superior observing conditions."
+ Lt should. therefore be replaced. with more accurate methods such as that of. Ixaiser. (19995). Rhodes. Relreeier Groth (2000). anc Ixuijken (1999). in future. more sensitive surveys.," It should therefore be replaced with more accurate methods such as that of Kaiser (1999b), Rhodes, Refregier Groth (2000), and Kuijken (1999), in future, more sensitive surveys."
+ We also used. our simulations to study. the cllect of seeing. exposure time ancl pixelisation on the sensitivity to 1e shear.," We also used our simulations to study the effect of seeing, exposure time and pixelisation on the sensitivity to the shear."
+ We found that increased seeing ENCIIM increases 10 noise almost linearly. with the primary loss being 1ο decreased number of usable galaxies.," We found that increased seeing FWHM increases the noise almost linearly, with the primary loss being the decreased number of usable galaxies."
+ In the seeing-lominated regime. the sensitivity to shear is only weakly lependent on exposure time.," In the seeing-dominated regime, the sensitivity to shear is only weakly dependent on exposure time."
+ As long as this regime holds. un is therefore more efficient. to tend towards wide rather jan deep weak lensing surveys.," As long as this regime holds, it is therefore more efficient to tend towards wide rather than deep weak lensing surveys."
+ Increased pixel scale hardly allects the sensitivity until the pixel. scale. is comparable to the seeing FWHAL at which point the method fails for typical grounc-basecl seeing.," Increased pixel scale hardly affects the sensitivity until the pixel scale is comparable to the seeing FWHM, at which point the method fails for typical ground-based seeing."
+ Thus. extreme oversampling of the PSE does not seem to be necessary.," Thus, extreme oversampling of the PSF does not seem to be necessary."
+ We also tested the claim by van Waerbeke et al., We also tested the claim by van Waerbeke et al.
+" that spurious shear signals on small scales (@zx 10"") could be sroclucecl by overlapping isophotes of neighboring galaxies.", that spurious shear signals on small scales $\theta \la 10''$ ) could be produced by overlapping isophotes of neighboring galaxies.
+ Using simulated images without input shear. we weakly detect this ellect on scales @Zi:Y.," Using simulated images without input shear, we weakly detect this effect on scales $\theta
+\la 1'$."
+ The rms amplitude of his effect is of about14... which is smaller but comparable o that expected for lensing.," The rms amplitude of this effect is of about, which is smaller but comparable to that expected for lensing."
+ Overlapping isophotes are thus ikelv to explain the excess power found on small scale by his group., Overlapping isophotes are thus likely to explain the excess power found on small scale by this group.
+ We would like to thank Thomas Erben. Peter Schneider. Yannick οπου. Roberto Alaoli ancl Auréllicn Thion for useful discussions.," We would like to thank Thomas Erben, Peter Schneider, Yannick Mellier, Roberto Maoli and Auréllien Thion for useful discussions."
+ We are indebtecl to Nick Ixaiser. for woviding us with the Lmeat software., We are indebted to Nick Kaiser for providing us with the Imcat software.
+" We acknowledge he invaluable use of URAL and. SlExtractor. during this research,", We acknowledge the invaluable use of IRAF and SExtractor during this research.
+ AR was supported by a Ελα postdoctoral ellowship from the EEC Lensing Network. and by a Wolfson College Research Fellowship.," AR was supported by a TMR postdoctoral fellowship from the EEC Lensing Network, and by a Wolfson College Research Fellowship."
+" DC acknowledges he ""Sonderforschungsbereich 375-95 fürr ophvssik"" der Deutschen LForsschunegsgecmmeinsschalt for inancial support."," DC acknowledges the “Sonderforschungsbereich 375-95 fürr sik"" der Deutschen schaft for financial support."
+" This work was supported by the TM ebvork cOravitational Lensing: New Constraints on Cosmology and the Distribution of Dark Matter"" of the EC under contract No.", This work was supported by the TMR Network “Gravitational Lensing: New Constraints on Cosmology and the Distribution of Dark Matter” of the EC under contract No.
+ ERBEAIRA-C'P97-0172., ERBFMRX-CT97-0172.
+2002).,.
+. These values are comparable to those for MSP W and provide further support for the presence of a shock in the [7 Tuc W system., These values are comparable to those for MSP W and provide further support for the presence of a shock in the 47 Tuc W system.
+ However. we enplasize that L7 Tuc Wis à fundamentally different binary than PSR DB1957|20 and similar eclipsing svstenis; such as £7 Tuc J. O. and R (Boedanov et al.," However, we emphasize that 47 Tuc W is a fundamentally different binary than PSR B1957+20 and similar eclipsing systems, such as 47 Tuc J, O, and R (Bogdanov et al."
+ 2005). which contain mach less massive (~0.03 M). ereatlv evolved secondary stars.," 2005), which contain much less massive $\sim$$0.03$ $_\odot$ ), greatly evolved secondary stars."
+ MSP W posseses mma characteristic similar to PSR JiT105310 (J1710). the only known MSP in the elobular cluster NGC 6397 (D'Anüicoetal.2001:Ferraroct2003:Sabbietal. 2003).," MSP W posseses many characteristic similar to PSR J1740–5340 (J1740), the only known MSP in the globular cluster NGC 6397 \citep{DAm01,Ferr03,Sabbi03}."
+" Specifically. J1710 exhibits radio eclipses. X-ray cluission which is secnunely haud and possibly variable (Grindlayctal.90059) and is bound toa ~0.3 AL. ""red strageler™ companion— that is currently filling its Roche lobe."," Specifically, J1740 exhibits radio eclipses, X-ray emission which is seemingly hard and possibly variable \citep{Grind02} and is bound to a $\sim$$0.3$ $_\odot$ “red straggler” companion that is currently filling its Roche lobe."
+ In addition. the shape of the IIa. line observed in this 32.5-hour binary svsteuni indicates the presence of a swept-back stream of gas protruding from the companion. as the one possibly present in [7 Tuc W (see Fig.," In addition, the shape of the $\alpha$ line observed in this 32.5-hour binary system indicates the presence of a swept-back stream of gas protruding from the companion, as the one possibly present in 47 Tuc W (see Fig."
+ 2)., 2).
+ However. uulike 17 Tuc W. J1710 shows no signature of wind-driven heating of the companion iu the photometric liebiteurve (Ikaluzux. 2003).," However, unlike 47 Tuc W, J1740 shows no signature of wind-driven heating of the companion in the photometric lightcurve \citep{Kal03,Oro03}."
+. This fiudiug is consistent with the recently revised value of £zδν10?! eres |1 (Bassa&Stappers which indicates that the fiux from the pulsar wind incident10.. ou the ~1.6 BR. companion is too stall to have a measurable effect.," This finding is consistent with the recently revised value of $\dot{E}\approx3.3\times10^{34}$ ergs $^{-1}$ \citep{Bassa04}, , which indicates that the flux from the pulsar wind incident on the $\sim$ 1.6 $_{\odot}$ companion is too small to have a measurable effect."
+ Therefore. in this svstem RL overflow is not the result of bloating by radiation from the MSP but rather a consequence of evolution off of the main sequence.," Therefore, in this system RL overflow is not the result of bloating by irradiation from the MSP but rather a consequence of evolution off of the main sequence."
+ Perhaps the most peculiar feature of the J171 binary is the presence of Ie I absorption localized im a thin longitudinal strip ucar the surface of the secondary (Ferraroetal.2003)., Perhaps the most peculiar feature of the J1740 binary is the presence of He I absorption localized in a thin longitudinal strip near the surface of the secondary \citep{Ferr03}.
+. The shape of such a heated region could be due to irradiation of the companion by a highly anisotropic pulsar wind. pretercutially emitted inu the orbital plane.," The shape of such a heated region could be due to irradiation of the companion by a highly anisotropic pulsar wind, preferentially emitted in the orbital plane."
+ However. this would require the MSP spin and orbital angular monieutuni vectors to be exactly aligned.," However, this would require the MSP spin and orbital angular momentum vectors to be exactly aligned."
+ The thin strip can be more plausibly explained by the existence of an equatorial wind. emianatiug from the tidally-locked companion star duc to its forced rapid co-rotatiou. that is interacting with the pulsar wind.," The thin strip can be more plausibly explained by the existence of an equatorial wind, emanating from the tidally-locked companion star due to its forced rapid co-rotation, that is interacting with the pulsar wind."
+ Alternatively. the Πο I absorption could originate in the swept back tail of material as the one shown in Figure 2 for I7 Tuc W. The unusual X-ray aud optical properties of 17 Tuc W can serve towards finally establishing the long-suspected connection between LAINBs and MSPs.," Alternatively, the He I absorption could originate in the swept back tail of material as the one shown in Figure 2 for 47 Tuc W. The unusual X-ray and optical properties of 47 Tuc W can serve towards finally establishing the long-suspected connection between LMXBs and MSPs."
+" This is now possible because the [£7 Tue W system appears to be more typical of a (LAINB system. as if coutaius a maiu-sequence companiou that mav be RL πιο, than a MSP binary svsten."," This is now possible because the 47 Tuc W system appears to be more typical of a qLMXB system, as it contains a main-sequence companion that may be RL filling, than a MSP binary system."
+" In addition. the N-rav. spectrum of 17 Tuc W svsteii exhibits remarkable similarities to that of the LAINB trausieut SAX JLlsds.l3658 (henceforth JISOS) in quiescence,"," In addition, the X-ray spectrum of 47 Tuc W system exhibits remarkable similarities to that of the LMXB transient SAX J1808.4–3658 (henceforth J1808) in quiescence."
+ observations of the latter system have revealed that during the long periods between outbursts (~2 wears). this 2.01-hour binary has an X-ray spectra which is sccminely purely uou- with DP—INE and an N-ray luminosity of Ey=5<10 cress bin the 0.5-10 keV baud (Campana 2002).," observations of the latter system have revealed that during the long periods between outbursts $\sim$ 2 years), this 2.01-hour binary has an X-ray spectrum which is seemingly purely non-thermal, with $\Gamma=1.4^{+0.6}_{-0.3}$ and an X-ray luminosity of $L_{X}=5\times10^{31}$ ergs $^{-1}$ in the 0.5-10 keV band \citep{Camp02}."
+. Moreover. recent optical obscrvatious of the bloated ~0.05 AL. brown diu companion of «1505 show sinusoidal variations at the orbital period (Campanaetal2001) muplviug a temperature difference of AT.ους4300 I& between the two faces of the conrpanion. quite simular to that reported here for 17 Tuc W. Finally. the observed non-thermal enüssion iu JLS08 also cannot account for the optical flux as well as radiating huuinosity (La2ενLO? eres 1) required to produce the observed leating of the secondary. star (seeFig.1ofCampanactal.2001).," Moreover, recent optical observations of the bloated $\sim$ 0.05 $_{\odot}$ brown dwarf companion of J1808 show sinusoidal variations at the orbital period \citep{Camp04}
+ implying a temperature difference of $\Delta T_{\rm c}
+\sim$$1000\pm300$ K between the two faces of the companion, quite similar to that reported here for 47 Tuc W. Finally, the observed non-thermal emission in J1808 also cannot account for the optical flux as well as irradiating luminosity $L_{\rm irr}\gtrsim4\times10^{33}$ ergs $^{-1}$ ) required to produce the observed heating of the secondary star \citep[see Fig. 1 of][]{Camp04}."
+. The X-ray aud optical properties of LF Tuc W and «1505 are compared in Table 2., The X-ray and optical properties of 47 Tuc W and J1808 are compared in Table 2.
+ The primary difference between JLSO8 iu quiescence and [7 Tuc Wis the lack of radio pulsations from the former., The primary difference between J1808 in quiescence and 47 Tuc W is the lack of radio pulsations from the former.
+ One possibility is that the radio beams are not favorably orieuted. reudering J1808 undetectable at radio waveleugths.," One possibility is that the radio beams are not favorably oriented, rendering J1808 undetectable at radio wavelengths."
+" Alternatively, the radio signals from the nascent MSP. expected in J1808 could be obscured by a diffuse circun-binary euvelope of material. formed by the gas flowing out from the companion that is ultimately expelled frou the system by the pulsar wind."," Alternatively, the radio signals from the nascent MSP expected in J1808 could be obscured by a diffuse circum-binary envelope of material, formed by the gas flowing out from the companion that is ultimately expelled from the system by the pulsar wind."
+ Such au euvelope is likely present around the 17 Tuc W system as well aud may explain. iu part. why this MSP is occulted for ~LO% of the orbit aud is rarely detected at radio frequencies (Freire.2005).," Such an envelope is likely present around the 47 Tuc W system as well and may explain, in part, why this MSP is occulted for $\sim$ of the orbit and is rarely detected at radio frequencies \citep{Freire05}."
+. If this is the case. it would be vet another conunon feature of the two svstenis.," If this is the case, it would be yet another common feature of the two systems."
+ We propose that the great similarities in the values of the power law iudex and the uou-thermal N-rav lunimositv derived for 17 Tuc W and JLS808 point to a conmmon origin of the non-thermal N-ray enussiou m both svsteis., We propose that the great similarities in the values of the power law index and the non-thermal X-ray luminosity derived for 47 Tuc W and J1808 point to a common origin of the non-thermal X-ray emission in both systems.
+ Also. it is hiehlv probable that the source of energv supplied to the heated companion iu both systems is the same aswell’.," Also, it is highly probable that the source of energy supplied to the heated companion in both systems is the same as."
+ If this is indeed the case. then. as originally sueecsted by Burderietal.(2003) and Campanaetal.(2001).. in the JLS08 system there exists a pulsar wind. powered bv the rotational energvof the NS. which is constraining the outflow of material from the companion aud preventing its accretion onto the NS.," If this is indeed the case, then, as originally suggested by \citet{Bur03} and \citet{Camp04}, in the J1808 system there exists a pulsar wind, powered by the rotational energyof the NS, which is constraining the outflow of material from the companion and preventing its accretion onto the NS."
+ As in the 17 Tuc W svstem. the prescuce of," As in the 47 Tuc W system, the presence of"
+This leaves us with the only possibility that the observed properties of BL Lacertae are the result of à transient short lasting phase.,This leaves us with the only possibility that the observed properties of BL Lacertae are the result of a transient short lasting phase.
+ We can envisage the following scenario. somewhat similar to that already suggested by Corbettetal. (1996).," We can envisage the following scenario, somewhat similar to that already suggested by \citet{cor96}."
+. BL Lacertae in its initial state has properties similar to the sources of the SDSS/NVSS sample., BL Lacertae in its initial state has properties similar to the sources of the SDSS/NVSS sample.
+ Indeed. these are massive early-type galaxies and a large number of them have narrow lines and radio luminosities similar to that of BL Lacertae.," Indeed, these are massive early-type galaxies and a large number of them have narrow lines and radio luminosities similar to that of BL Lacertae."
+ From the point of view of their optical spectra they are LEG and lack broad lines., From the point of view of their optical spectra they are LEG and lack broad lines.
+ Subsequently (possibly ~20 years ago). its BLR underwent an increase of luminosity due to an increased amount of cold gas in the nuclear regions and/or to a higher level of ionizing continuum.," Subsequently (possibly $\sim 20$ years ago), its BLR underwent an increase of luminosity due to an increased amount of cold gas in the nuclear regions and/or to a higher level of ionizing continuum."
+ These two effects may even be related and caused by a fresh input of accreting gas., These two effects may even be related and caused by a fresh input of accreting gas.
+ The BLR structure might not have yet reached a stable configuration. accounting for its different properties when compared to other AGN.," The BLR structure might not have yet reached a stable configuration, accounting for its different properties when compared to other AGN."
+" Also the NLR luminosity will grow with time and will also eventually change its state of ionization, but on a much larger timescale with respect to the BLR."," Also the NLR luminosity will grow with time and will also eventually change its state of ionization, but on a much larger timescale with respect to the BLR."
+ Based on the analysis of a single object it ts clearly impossible to set a timescale for the duration of the putative bright phase., Based on the analysis of a single object it is clearly impossible to set a timescale for the duration of the putative bright phase.
+ Furthermore. BL Lacertae was probably discovered since this object has been subject to repeated spectroscopic observations.," Furthermore, BL Lacertae was probably discovered since this object has been subject to repeated spectroscopic observations."
+ However. our failure to find objects in the local Universe that might constitute its parent population suggests that the timescale associated with the period of high accretion must be orders of magnitude shorter than the lifetime of radio-loud AGN.," However, our failure to find objects in the local Universe that might constitute its parent population suggests that the timescale associated with the period of high accretion must be orders of magnitude shorter than the lifetime of radio-loud AGN."
+ An alternative possibility is that the birth of the BLR marks the transition from a low-power radio galaxy to a high-power source., An alternative possibility is that the birth of the BLR marks the transition from a low-power radio galaxy to a high-power source.
+ This would require a rapid increase in the luminosity of the large-scale radio structures to reach the level observed in tthe HEG of the 3CR sample. within a sufficiently short time so as not to produce a substantial population of transient sources.," This would require a rapid increase in the luminosity of the large-scale radio structures to reach the level observed in the HEG of the 3CR sample, within a sufficiently short time so as not to produce a substantial population of transient sources."
+ Instead. the available data rule out that BL Lacertae became an AGN only very recently. tthat we are witnessing its birth. because its radio emission extends ~10 kpe away from the core.," Instead, the available data rule out that BL Lacertae became an AGN only very recently, that we are witnessing its birth, because its radio emission extends $\sim 10$ kpc away from the core."
+ This implies that this source is active since at least ~3x10° years. assuming an expansion speed of 0.1 c. We conclude that the parent population of BL Lacertae can be found among the large population of miniature radio-loud AGN forming the SDSS/NVSS sample. but this also requires that this object is experiencing a short transient phase.," This implies that this source is active since at least $\sim 3 \times 10^5$ years, assuming an expansion speed of 0.1 c. We conclude that the parent population of BL Lacertae can be found among the large population of miniature radio-loud AGN forming the SDSS/NVSS sample, but this also requires that this object is experiencing a short transient phase."
+ A continuation of the spectroscopic monitoring of this peculiar source caught in a crucial phase of its evolution can help us tremendously in our study of the physics and evolution of these systems., A continuation of the spectroscopic monitoring of this peculiar source caught in a crucial phase of its evolution can help us tremendously in our study of the physics and evolution of these systems.
+"of Lyman-limit systems of Pérouxetal.(2001),, where they report y=2.4573 05. ","of Lyman-limit systems of \citet{Per01}, where they report $\gamma=2.45^{+0.75}_{-0.04}$ ."
+"From z=4.5 to z=3, we see a decrease in the abundance by a factor of about two in both simulations and the observations, and the agreement between the two is very good, although the simulation points tend to fall slightly below the observations."," From $z=4.5$ to $z=3$, we see a decrease in the abundance by a factor of about two in both simulations and the observations, and the agreement between the two is very good, although the simulation points tend to fall slightly below the observations."
+" From z—3 to z—1, the simulation (D5) suggests a further rapid decrease in DLA abundance by a factor of ~6, which is not seen at this level in the existing observations."," From $z=3$ to $z=1$, the simulation (D5) suggests a further rapid decrease in DLA abundance by a factor of $\sim 6$, which is not seen at this level in the existing observations."
+" But the observational data at low redshift are still relatively uncertain, as indicated by the large error bars."," But the observational data at low redshift are still relatively uncertain, as indicated by the large error bars."
+" The rapid decline is also reflected in the fact that Ώηι decreases from 0.66 (z= 3) to 0.14 (z 1) in D5 over this redshift range, a reduction of nearly a factor of 5 (see Figure 1))."," The rapid decline is also reflected in the fact that $\OHI$ decreases from 0.66 $z=3$ ) to 0.14 $z=1$ ) in D5 over this redshift range, a reduction of nearly a factor of 5 (see Figure \ref{omega.eps}) )."
+" If this significant decrease in the number of DLAs from z=3 to z=1 is real, it would partly explain why it is so difficult to find DLAs at z<1."," If this significant decrease in the number of DLAs from $z=3$ to $z=1$ is real, it would partly explain why it is so difficult to find DLAs at $z\leq 1$."
+" On the other hand, not much evolution is seen from z =1to z=0 in the G5 simulation."," On the other hand, not much evolution is seen from $z=1$ to $z=0$ in the G5 simulation."
+ This is related to the fact that Qur in G5 does not decrease very much from z=1 (Qni= 0.19) to z=0 (Ώπι= 0.16)., This is related to the fact that $\OHI$ in G5 does not decrease very much from $z=1$ $\OHI=0.19$ ) to $z=0$ $\OHI=0.16$ ).
+" However, as discussed earlier, our power-law fits to the opia—Mt relation are not well constrained for z=0 (and possibly for z=1 as well), so the results at z<2 should be interpreted with caution."," However, as discussed earlier, our power-law fits to the $\sdla - {M_{\rm tot}}$ relation are not well constrained for $z=0$ (and possibly for $z=1$ as well), so the results at $z\leq 2$ should be interpreted with caution."
+" At z>3, we saw that lower resolution runs tend to predict a larger abundance due to a shallower slope in the relation between the DLA cross-section and the halo mass, but it is not clear if other forms of systematic bias dominate at very low redshift for simulations with poor resolution."," At $z\geq 3$, we saw that lower resolution runs tend to predict a larger abundance due to a shallower slope in the relation between the DLA cross-section and the halo mass, but it is not clear if other forms of systematic bias dominate at very low redshift for simulations with poor resolution."
+" We will need yet higher resolution simulations with large box-sizes to make a more robust prediction of the DLA abundance at zX2, and until then, it is not clear whether the current results for DLA abundance at z<2, which tend to fall below the observational data, are trustworthy."," We will need yet higher resolution simulations with large box-sizes to make a more robust prediction of the DLA abundance at $z\leq 2$, and until then, it is not clear whether the current results for DLA abundance at $z\leq 2$, which tend to fall below the observational data, are trustworthy."
+ This is why we have widened the shaded confidence region in Figure 6 significantly for z«2., This is why we have widened the shaded confidence region in Figure \ref{cum_evolve.eps} significantly for $z\leq 2$.
+" The column density distribution function f(N,X(z)) is defined such that f(N,X)aNd.X is the number of absorbers per sight line with column densities in the interval [N,N+dN], and absorption distances in the interval dX]."," The column density distribution function $f(N,X(z))$ is defined such that $f(N,X){\rm d}N{\rm d}X$ is the number of absorbers per sight line with column densities in the interval $[N,N+{\rm d}N]$, and absorption distances in the interval $[X,X+{\rm d}X]$ ."
+" The absorption distance X(z) is given by X(z) = / δε» This definition is based on an argument by Bahcall&Peebles (1969),, who pointed out that the probability of absorption for a quasar sight-line in the redshift interval [z,z4-dz] is dPος(1+z)?dr(14z)?[Ho/H(z)]dz=dX."," The absorption distance $X(z)$ is given by X(z) = _0^z (1+z')^2 This definition is based on an argument by \citet{Bah69}, , who pointed out that the probability of absorption for a quasar sight-line in the redshift interval $[z,z+{\rm
+d}z]$ is ${\rm d}P\propto (1+z)^2 {\rm d}r \propto (1+z)^2 [H_0 /H(z)]
+{\rm d}z \equiv {\rm d}X$."
+" In practice, if the comoving box-size of the simulation is AL, then the corresponding absorption distance per sight-line is AX=(Ho/c)(1+z)°AL."," In practice, if the comoving box-size of the simulation is $\Del L$, then the corresponding absorption distance per sight-line is $\Del X = ({H_0}/{c})(1+z)^2 \Del L$."
+" For example, for AL=10h~'Mpc and z=3, we have AX= 0.0534."," For example, for $\Del L=10\himpc$ and $z=3$, we have $\Del X= 0.0534$ ."
+" Assuming that DLAs do not overlap along a sight-through the simulation volume (which is a very good approximation given the small size of the simulation box, where the expected number of DLAs per sight-line at z=3 for a 10h~'Mpc pathis zz 10”3), we can compute theNur distribution function by counting the number of grid-cells"," Assuming that DLAs do not overlap along a sight-linethrough the simulation volume (which is a very good approximation given the small size of the simulation box, where the expected number of DLAs per sight-line at $z=3$ for a $10\himpc$ pathis $\approx 10^{-3}$ ), we can compute the$\NHI$ distribution function by counting the number of grid-cells"
+Polarization could only be studied for the second epoch and using intra-European baselines.,Polarization could only be studied for the second epoch and using intra-European baselines.
+ reff.pol provides a polarization map for this epoch., \\ref{f.pol} provides a polarization map for this epoch.
+ The mean fractional polarization ts (peak of 1.3%)). which is barely in excess of the one reported from MOJAVE at 15 GHz at the same epoch.," The mean fractional polarization is (peak of ), which is barely in excess of the one reported from MOJAVE at 15 GHz at the same epoch."
+ The absolute polarization angle cannot be determined for our data., The absolute polarization angle cannot be determined for our data.
+ To our knowledge. this is the first report of à VLBI polarization experiment using the new 40m Yebes radio telescope.," To our knowledge, this is the first report of a VLBI polarization experiment using the new 40m Yebes radio telescope."
+ The census of extragalactic gamma-ray sources throughout the EGRET era has been dominated by blazars (flat spectrum radio quasars and BL Lac type objects)., The census of extragalactic gamma-ray sources throughout the EGRET era has been dominated by blazars (flat spectrum radio quasars and BL Lac type objects).
+ The discovery of gamma-ray emission from radio-loud narrow-line Seyfert! nuclei by is therefore of great importance to the study of relativistic jets in AGNs (Abdoetal.2009d)., The discovery of gamma-ray emission from radio-loud narrow-line Seyfert1 nuclei by is therefore of great importance to the study of relativistic jets in AGNs \citep{sample}.
+. First. it has confirmed that these relativistic structures exist in radio loud LSI. as proposed by Zhouetal.(2003) and Doietal. (2006).," First, it has confirmed that these relativistic structures exist in radio loud NLS1, as proposed by \citet{Zhou2003} and \citet{Doi2006}."
+. Moreover. exploiting the sensitivity and the surveying capability ofFermi. it has triggered the organization of the large AWL campaign including the global e-VLBI observations presented here (Abdoetal.2009c).," Moreover, exploiting the sensitivity and the surveying capability of, it has triggered the organization of the large MWL campaign including the global e-VLBI observations presented here \citep{mwl}."
+. The MWL campaign has revealed variability at all energies. and permitted the estimate of the physical parameters of the jet. such as the dissipation radius (p=6.7x10'? cem). the magnetic field (B=4.1 GG). and the jet total power (Abdoetal. 2009c).," The MWL campaign has revealed variability at all energies, and permitted the estimate of the physical parameters of the jet, such as the dissipation radius $r=6.7\times 10^{16}$ cm), the magnetic field $B=4.1$ G), and the jet total power \citep{mwl}."
+. A bulk Lorentz factor [>1 was implied by the large Compton dominance: in particular. Γ510 was assumed as a typical value.," A bulk Lorentz factor $\Gamma>1$ was implied by the large Compton dominance; in particular, $\Gamma=10$ was assumed as a typical value."
+ Our lower limit to the brightness temperatures estimated by elliptical Gaussian model fitting to the global e-VLBI visibilities is a few »10!! K and provides an independent confirmation that this jet has a Doppler factor greater than one (seealsoDotetal.2006:Zhou2003).. as well as in other RL-NLSI (Gu&Chen2010).," Our lower limit to the brightness temperatures estimated by elliptical Gaussian model fitting to the global e-VLBI visibilities is a few $\times 10^{11}$ K and provides an independent confirmation that this jet has a Doppler factor greater than one \citep[see also][]{Doi2006,Zhou2003}, as well as in other RL-NLS1 \citep{Gu2010}."
+. Because of the uncertainty in the overall amplitude scaling. a Doppler factor estimate based on the variability brightness temperature Is not discussed in this study.," Because of the uncertainty in the overall amplitude scaling, a Doppler factor estimate based on the variability brightness temperature is not discussed in this study."
+ However. a trend in the total flux density from our observation is present. with a peak measured on 2009 May 23 and a decrease in the following epochs. which are also characteristic features of blazar-like emission.," However, a trend in the total flux density from our observation is present, with a peak measured on 2009 May 23 and a decrease in the following epochs, which are also characteristic features of blazar-like emission."
+ This result also agrees with other independent variability measurements requiring 0>| reported at various frequencies in Abdoetal. (2009¢)., This result also agrees with other independent variability measurements requiring $\delta>1$ reported at various frequencies in \citet{mwl}.
+. Significant fractional polarization is also a characteristic of detected Jets (Hovattaetal.2010)., Significant fractional polarization is also a characteristic of detected jets \citep{Hovatta2010}.
+. In JO948+0022 itself. Foschinietal.(2010b) reported an increase in the fractional polarization up to and a swing of the EVPA of about 90° in possible connection with the 2010 gamma-ray outburst.," In J0948+0022 itself, \citet{Foschini2011} reported an increase in the fractional polarization up to and a swing of the EVPA of about $90^\circ$ in possible connection with the 2010 gamma-ray outburst."
+ This is similar to what has been reported for other blazars. such as 115024106 (Abdoetal.2010).. and might be related to the physical mechanisms responsible for enhanced gamma-ray activity.," This is similar to what has been reported for other blazars, such as 1502+106 \citep{Abdo2010}, and might be related to the physical mechanisms responsible for enhanced gamma-ray activity."
+ In this context. our detected value of ~1% fractional polarization seems more characteristic of an intermediate state of activity. consistent with the transition to a low state at high energy observed towards the end of the 2009 MWL campaign (Abdoetal.2009€).," In this context, our detected value of $\sim 1\%$ fractional polarization seems more characteristic of an intermediate state of activity, consistent with the transition to a low state at high energy observed towards the end of the 2009 MWL campaign \citep{mwl}."
+. In. this paper. we have presented radio. interferometric observations in. real-time. e-VLBI mode (Szomoru2008). using. for the first time. a global array spanning three continents.," In this paper, we have presented radio interferometric observations in real-time e-VLBI mode \citep{Szomoru2008}, using, for the first time, a global array spanning three continents."
+ The feasibility of transferring data across continents and correlating them in real time has been demonstrated in a number of tests and in particular in the LBA observations of 11987A. which were correlated at JIVE (Tingayetal.2009:etal.2011).," The feasibility of transferring data across continents and correlating them in real time has been demonstrated in a number of tests and in particular in the LBA observations of 1987A, which were correlated at JIVE \citep{Tingay2009,Ng2011}."
+. However. these are the first astronomical science observations on a global scale.," However, these are the first astronomical e-VLBI science observations on a global scale."
+ From the very beginning of VLBI. global baselines have been routinely achieved between e.g.. Europe and the US. and some examples exist of truly global collaborations between arrays on three different continents (seee.g.Fomalontetal.2001).," From the very beginning of VLBI, global baselines have been routinely achieved between e.g., Europe and the US, and some examples exist of truly global collaborations between arrays on three different continents \citep[see
+ e.g.][]{Fomalont2001}."
+. However. it has always been a challenge to organise these experiments. especially in the case of time-critical observations such as our coordinated MWL campaign.," However, it has always been a challenge to organise these experiments, especially in the case of time-critical observations such as our coordinated MWL campaign."
+ The advantages of performing such observations in real time e-VLBI (when transporting the data through high-speed networks is possible. from all participating telescopes) include the ability to monitor the network performance in real-time and the possibility of delivering prompt results (Girolettietal. 2010a.b).," The advantages of performing such observations in real time e-VLBI (when transporting the data through high-speed networks is possible from all participating telescopes) include the ability to monitor the network performance in real-time and the possibility of delivering prompt results \citep{Giroletti2010a,Giroletti2010b}."
+. But most importantly. with the standardization of the data formats and transport protocols (Whitneyetal.2009).. it will become much easier to create ad-hoc global networks at relatively short notice. providing excellent wv-coverage and unique long baselines. as well as greatly increased sensitivity due to the increase in sustainable data rates and the number of telescopes.," But most importantly, with the standardization of the data formats and transport protocols \citep{Whitney2009}, it will become much easier to create ad-hoc global networks at relatively short notice, providing excellent $uv$ -coverage and unique long baselines, as well as greatly increased sensitivity due to the increase in sustainable data rates and the number of telescopes."
+ Our pioneering global e-VLBI efforts. in spite of some initial technical difficulties. have convincingly demonstrated the operational feasibility of this type of observation.," Our pioneering global e-VLBI efforts, in spite of some initial technical difficulties, have convincingly demonstrated the operational feasibility of this type of observation."
+ Using global baselines and a high observing frequency. we have probed the innermost region at or above the self-absorption frequency. thus obtaining the most valuable information às far as the comparison to high energy activity is concerned.," Using global baselines and a high observing frequency, we have probed the innermost region at or above the self-absorption frequency, thus obtaining the most valuable information as far as the comparison to high energy activity is concerned."
+ The overall success of the MWL campaign and in particular of the global e-VLBI observations presented here is highly encouraging for the continuation of the synergy between high energy astrophysics and VLBI., The overall success of the MWL campaign and in particular of the global e-VLBI observations presented here is highly encouraging for the continuation of the synergy between high energy astrophysics and VLBI.
+with their intense magnetic fields. possess significant cllipticities. making them good candidates for gravitational vVAVO ποιος (Bonazzola&Gourgoulhon1996:AlelatosPayne2005:Stellactal.Laskellοἱ2008:Dall'Ossoe2009).,"with their intense magnetic fields, possess significant ellipticities, making them good candidates for gravitational wave sources \citep{bg96,mp05,setal05,hetal08,detal09}."
+". Recent data from the fifth. Laser Interferometer Gravitational Wave Observatory (LIGO) Science Run set an upper limi ofcL4.10! on the Crab Pulsar (Abbottetal.2008.2010)... translating into an internal magnetic [eld of 1072 ""p under standard assumptions."," Recent data from the fifth Laser Interferometer Gravitational Wave Observatory (LIGO) Science Run set an upper limit of $\epsilon\lesssim 1.4\times 10^{-4}$ on the Crab Pulsar \citep{aetal08,aetal10}, translating into an internal magnetic field of $\lesssim 10^{12}$ T under standard assumptions."
+ LIGO non-detections of the central compact object (CXCO) in the supernova remnan Cassiopeia A (Cas A) have constrained. its ellipticity as well., LIGO non-detections of the central compact object (CCO) in the supernova remnant Cassiopeia A (Cas A) have constrained its ellipticity as well.
+ The Cas A οςο has not been detected. clectromagnetically. making i impossible to infer its external magnetic field. from the spin-down rate.," The Cas A CCO has not been detected electromagnetically, making it impossible to infer its external magnetic field from the spin-down rate."
+" However. Wetteetal.(2008). and. Wette.(2010) constrained its ellipticity as a function of gravitational wave [requeney (eg. cE3.6103 for 100 Liz. ος0.6)1 for 20t Hz. and ος0.88.107 for 300 Lz). implying an internal magnetic field 107 7E. Lastly. Chungetal.(2011) gaj0wed tha it will be possible to use future data from LIGO to set a lower limit of LO"" T on the magnetic lield (e101) of jo putative 24-vear-old. neutron star in the supernova remnant SNR I987A. Ciravitational waves are generated by a rotating star when it is not spherically svmametric and when its “wohe angle! 8. the angle between its total angular momentum vector and symmetry axis. is nonzero."," However, \citet{wetal08} and \citet{w10} constrained its ellipticity as a function of gravitational wave frequency (e.g., $\epsilon\lesssim 3.6 \times 10^{-4}$ for 100 Hz, $\epsilon\lesssim 0.6\times 10^{-4}$ for 200 Hz, and $\epsilon\lesssim 0.38\times 10^{-4}$ for 300 Hz), implying an internal magnetic field $\lesssim 10^{14}$ T. Lastly, \citet{cetal11} showed that it will be possible to use future data from LIGO to set a lower limit of $10^7$ T on the magnetic field $\epsilon \lesssim 10^{-4}$ ) of the putative 24-year-old neutron star in the supernova remnant SNR 1987A. Gravitational waves are generated by a rotating star when it is not spherically symmetric and when its `wobble angle' $\theta$, the angle between its total angular momentum vector and symmetry axis, is nonzero."
+ In eeneral. therefore. the star precesses as it radiates.," In general, therefore, the star precesses as it radiates."
+ The magnetic field analysed in this paper is axisvmametric. which deforms the star into an οἱipsoid.," The magnetic field analysed in this paper is axisymmetric, which deforms the star into an ellipsoid."
+ The most general expression for the gravitational wave signal of a precessing triaxial ellipsoid given by. e.g.. Eqs. (," The most general expression for the gravitational wave signal of a precessing triaxial ellipsoid [given by, e.g., Eqs. ("
+19)(26) of (Jaranowski.Ixrólak.&Schutz 20013]] depends on 8 and the angle + between the angular momentum vector and line of sight to the Earth. as well as e.,"19)–(26) of \citep{jks98}] ] depends on $\theta$ and the angle $\iota$ between the angular momentum vector and line of sight to the Earth, as well as $\epsilon$."
+ The signal is strongest when 6=z/2 (ic. the rotation axis is perpendicular to the symmetry axis. which is in turn parallel to the magnetic axis) and 7=0 (i.e. the rotation axis is directed towards Earth).," The signal is strongest when $\theta=\pi/2$ (i.e., the rotation axis is perpendicular to the symmetry axis, which is in turn parallel to the magnetic axis) and $\iota=0$ (i.e., the rotation axis is directed towards Earth)."
+ In ttjs paper. to simplify matters. we henceforth assume implicitly that 6=x/2 ands =0. so that gravitational wave emission ancl signal detection are assumed to be optimal and there is no precession. in order to focus on ο without geometric complicalions.," In this paper, to simplify matters, we henceforth assume implicitly that $\theta=\pi/2$ and $\iota=0$, so that gravitational wave emission and signal detection are assumed to be optimal and there is no precession, in order to focus on $\epsilon$ without geometric complications."
+ In Mastranoctal.(2011)... we constructed hyvdromagnetie equilibria for stratified. stars.," In \citet{metal11}, we constructed hydromagnetic equilibria for stratified, stars."
+" Thο ConinirorA adopted barotropic assumption. while simplifving calculations. severely restricts the form of the field that can be ""fitted into he star. e.g. Haskelletal.(2008). found that the field must. vanish at the surface. contrary to observations. ancl and Ciolfi.Ferrari.&Cualtieri(2010) found that only configurations dominated. by the poloidal component (poloidal energy ο90% of total) are allowed. contrary to the numerical simulations of Braithwaite&Nordlund(2006)."," The commonly adopted barotropic assumption, while simplifying calculations, severely restricts the form of the field that can be `fitted' into the star, e.g., \citet{hetal08} found that the field must vanish at the surface, contrary to observations, and \citet{lj09} and \citet{cfg10} found that only configurations dominated by the poloidal component (poloidal energy $\gtrsim 90\%$ of total) are allowed, contrary to the numerical simulations of \citet{bn06}."
+. By abandoning the barotropic assumption. we are able to construct a simple. self-consistent hvdromagnetie equilibrium with an internal field that can be matched to an external dipole.," By abandoning the barotropic assumption, we are able to construct a simple, self-consistent hydromagnetic equilibrium with an internal field that can be matched to an external dipole."
+" In our configuration. because we do not require the pressure to be a ""unction of density alone. we are less restricted in the choices of poloidal and toroidal components: they need not be of any xuwticular relative strengths ancl are independently adjustable."," In our configuration, because we do not require the pressure to be a function of density alone, we are less restricted in the choices of poloidal and toroidal components; they need not be of any particular relative strengths and are independently adjustable."
+ Locidentally. this means that the magnetic field configuration is independent of the equation of state chosen. cf.," Incidentally, this means that the magnetic field configuration is independent of the equation of state chosen, cf."
+ Lander&Jones(2009) and Ciolfi.Ferrari.&Cualtieri(2010)., \citet{lj09} and \citet{cfg10}.
+. In this paper. we present one possible astrophysical application of the aforementioned. result. namely to constrain the internal magnetic fields of neutron stars in conjunction with gravitational wave observations.," In this paper, we present one possible astrophysical application of the aforementioned result, namely to constrain the internal magnetic fields of neutron stars in conjunction with gravitational wave observations."
+ As we can match the internal field to an external dipole. we are able to relate one set of observations (external magnetic field strength. from spin period and spin-down rate) with another (gravitational wave upper limits). at least in principle.," As we can match the internal field to an external dipole, we are able to relate one set of observations (external magnetic field strength, from spin period and spin-down rate) with another (gravitational wave upper limits), at least in principle."
+ In Sec., In Sec.
+ 2. we brielly. describe our field structure and summarize the results of Mastranoetal.(2011).," 2, we briefly describe our field structure and summarize the results of \citet{metal11}."
+.. Phen. in Sec.," Then, in Sec."
+ 3. we generalize the earlier wor‘k to include a superconducting interior.," 3, we generalize the earlier work to include a superconducting interior."
+ We compare the cllipticity calculated using our superconducting mocdel to the gravitzuional-wave upper limits of SL known millisecond pulsars. including also the important ellect of accretion-induced. diamagnetic shielding.," We compare the ellipticity calculated using our superconducting model to the gravitational-wave upper limits of 81 known millisecond pulsars, including also the important effect of accretion-induced diamagnetic shielding."
+ Lastly. in Sec.," Lastly, in Sec."
+ 4. we summarize our results and discuss the possibilitythat the internal fields of millisecond pulsars are stronger than currently thought.," 4, we summarize our results and discuss the possibilitythat the internal fields of millisecond pulsars are stronger than currently thought."
+ Alastranoetal.(2011). considered. a general class of poloidal-toroidal magnetic field. configurations. which are broadly representative of the field structures observed in numerical simulations and exhibit the following properties: We write the magnetic field in the form pioneered by Chandrasekhar (19," \citet{metal11} considered a general class of poloidal-toroidal magnetic field configurations, which are broadly representative of the field structures observed in numerical simulations \citep{bn06,bs06,b09} and exhibit the following properties: We write the magnetic field in the form pioneered by \citet{c56}, , +"
+ Alastranoetal.(2011). considered. a general class of poloidal-toroidal magnetic field. configurations. which are broadly representative of the field structures observed in numerical simulations and exhibit the following properties: We write the magnetic field in the form pioneered by Chandrasekhar (195," \citet{metal11} considered a general class of poloidal-toroidal magnetic field configurations, which are broadly representative of the field structures observed in numerical simulations \citep{bn06,bs06,b09} and exhibit the following properties: We write the magnetic field in the form pioneered by \citet{c56}, , +"
+ Alastranoetal.(2011). considered. a general class of poloidal-toroidal magnetic field. configurations. which are broadly representative of the field structures observed in numerical simulations and exhibit the following properties: We write the magnetic field in the form pioneered by Chandrasekhar (1956," \citet{metal11} considered a general class of poloidal-toroidal magnetic field configurations, which are broadly representative of the field structures observed in numerical simulations \citep{bn06,bs06,b09} and exhibit the following properties: We write the magnetic field in the form pioneered by \citet{c56}, , +"
+ Alastranoetal.(2011). considered. a general class of poloidal-toroidal magnetic field. configurations. which are broadly representative of the field structures observed in numerical simulations and exhibit the following properties: We write the magnetic field in the form pioneered by Chandrasekhar (1956)," \citet{metal11} considered a general class of poloidal-toroidal magnetic field configurations, which are broadly representative of the field structures observed in numerical simulations \citep{bn06,bs06,b09} and exhibit the following properties: We write the magnetic field in the form pioneered by \citet{c56}, , +"
+ Alastranoetal.(2011). considered. a general class of poloidal-toroidal magnetic field. configurations. which are broadly representative of the field structures observed in numerical simulations and exhibit the following properties: We write the magnetic field in the form pioneered by Chandrasekhar (1956).," \citet{metal11} considered a general class of poloidal-toroidal magnetic field configurations, which are broadly representative of the field structures observed in numerical simulations \citep{bn06,bs06,b09} and exhibit the following properties: We write the magnetic field in the form pioneered by \citet{c56}, , +"
+ Alastranoetal.(2011). considered. a general class of poloidal-toroidal magnetic field. configurations. which are broadly representative of the field structures observed in numerical simulations and exhibit the following properties: We write the magnetic field in the form pioneered by Chandrasekhar (1956)..," \citet{metal11} considered a general class of poloidal-toroidal magnetic field configurations, which are broadly representative of the field structures observed in numerical simulations \citep{bn06,bs06,b09} and exhibit the following properties: We write the magnetic field in the form pioneered by \citet{c56}, , +"
+We investigate low the ( [actor is allected if OPD is deviated from the optimal value aud iud that a deviated OPD (οιain) does not result in a significalu Q factor degradation. which is uitigated as 2 increases.,"We investigate how the $Q$ factor is affected if OPD is deviated from the optimal value and find that a deviated OPD (5mm) does not result in a significant $Q$ factor degradation, which is mitigated as $R$ increases."
+ We iud that the Q [actor increases \albh FR for both DEDI aud DE. aud eventually converge at very high A (2> 100.000).," We find that the $Q$ factor increases with $R$ for both DFDI and DE, and eventually converge at very high $R$ $R\geq$ 100,000)."
+ The couvergerce of DEDI aud DE inethods is a iatural Consequence because tje Measurement method does not 1ake a difference alter the spectral 'esolutiou becomes extremely high., The convergence of DFDI and DE methods is a natural consequence because the measurement method does not make a difference after the spectral resolution becomes extremely high.
+ In addition. Q factors at a given £e increase as Tay drops from 91001 to 2£00Ix. which is due ο stronger molecular absorption features in NIB (see Fie. 3)).," In addition, $Q$ factors at a given $R$ increase as $T_{\rm{eff}}$ drops from 3100K to 2400K, which is due to stronger molecular absorption features in NIR (see Fig. \ref{fig:Wav_Flux}) )."
+ The Q actor decreases as Vsini increases because stellar rotation broacdeus the absorption lines. leacling o less seusitive Iueasureiment.," The $Q$ factor decreases as $V \sin{i}$ increases because stellar rotation broadens the absorption lines, leading to less sensitive measurement."
+ We compare 6) [actors lor both DEDI aud DE at a given. Zt., We compare $Q$ factors for both DFDI and DE at a given $R$ .
+" For slow rotators η, DFDI is nore advantageous over DE at low aud medium 2 (5.000 to 20.000) lor the same wavelength coverage AA."," For slow rotators $0\ \rm{km\cdot s}^{-1}\le V \sin{i}\le2\ \rm{km\cdot s}^{-1}$ ), DFDI is more advantageous over DE at low and medium $R$ (5,000 to 20,000) for the same wavelength coverage $\Delta\lambda$."
+ The iniproveieut Of DEDI compared to DE is ~3.1 (2=5.000). 7.2.| (R=10.000) and ~1.7 (2R=20.000). respectively.," The improvement of DFDI compared to DE is $\sim$ 3.1 $R$ =5,000), $\sim$ 2.4 $R$ =10,000) and $\sim$ 1.7 $R$ =20,000), respectively."
+ In other words. optimized DEDI with £2 of 2.000. 10.000 aud 20.000 are equivalent in Dopper seusitivity to DE with @ of 16.000. 2L000 ancl 3LO000. respectively.," In other words, optimized DFDI with $R$ of 5,000, 10,000 and 20,000 are equivalent in Doppler sensitivity to DE with $R$ of 16,000, 24,000 and 34,000, respectively."
+ Tle provement of DEDI at [t 20.000 to 50.000 1s not as noticeable as at low # range.," The improvement of DFDI at $R$ 20,000 to 50,000 is not as noticeable as at low $R$ range."
+ The differeuce between DEDI anc DE becomes negligible when /? is over 100.000.," The difference between DFDI and DE becomes negligible when $R$ is over 100,000."
+ For ‘elatively faster roators (5ανν€VsinicUI0kin-s 1). the provement with DEDI is less ovious than it is for very slow rotaors.," For relatively faster rotators $5\ \rm{km\cdot s}^{-1}\le V \sin{i}\le10\ \rm{km\cdot s}^{-1}$ ), the improvement with DFDI is less obvious than it is for very slow rotators."
+ DEDI has strength when tlie spectral lines in a stellar spectrum afe HOl resoved by a spectrograph. which is tve case for low aud medium resolution specgraph.," DFDI has strength when the spectral lines in a stellar spectrum are not resolved by a spectrograph, which is the case for low and medium resolution spectrograph."
+ Under sich. coucitious. the fixed delay inte‘ferometer provides additional resolving powers for tje system.," Under such conditions, the fixed delay interferometer provides additional resolving powers for the system."
+ After the lines are Fully resolved by tle spectrograph itsell. tle interferonjeter in tle systeu becomes dispeusable. which is the reason wvo Wwe see tlie convergelce of DEDI aud DE a very lugh spectral resolution.," After the lines are fully resolved by the spectrograph itself, the interferometer in the system becomes dispensable, which is the reason why we see the convergence of DFDI and DE at very high spectral resolution."
+ Fundajeutal. performaice of a Fourler-transform spectremeter (FTS) in tle applicatiou of Doppler measurements Las een discussed by. 2.., Fundamental performance of a Fourier-transform spectrometer (FTS) in the application of Doppler measurements has been discussed by \citet{Maillard1996}.
+ There are similarities between the FTS aud tl DFDI method. for example: 1. both methods 1se the interferometer as a fine s»ectral resolvi element: 2. RV is 1neasurec by monitoring the temporal phase change at a fixed OPD of tl interferometer.," There are similarities between the FTS and the DFDI method, for example: 1, both methods use the interferometer as a fine spectral resolving element; 2, RV is measured by monitoring the temporal phase change at a fixed OPD of the interferometer."
+ I DEDI met1ος. OPD is scaunec in each frequency channel because of two relativel: tiltecl inirrors. aid the resolution of the post-disperser in DEDI is chosen to ensure a reasotab (tinge visibility.," In DFDI method, OPD is scanned in each frequency channel because of two relatively tilted mirrors, and the resolution of the post-disperser in DFDI is chosen to ensure a reasonable fringe visibility."
+ Therefore. he DEDI method is a exteudec version of the FTS method with a low-inecil Jesoution post-disperser.," Therefore, the DFDI method is a extended version of the FTS method with a low-medium resolution post-disperser."
+ However. one major clierence between these two met10€S is that iiterferoimeter itself is used as a spectrometer by OPD scanuing in the FETS mehod while ai (itional spectrograph is employed iu the DEDI mehod.," However, one major difference between these two methods is that the interferometer itself is used as a spectrometer by OPD scanning in the FTS method while an additional spectrograph is employed in the DFDI method."
+ The advantage of introdicing an additia| spectrograph into the system is that the visibility (or fringe contrast) is uo loüuger limitect 1 bandpass as in the FTS case. which is the reason that the DEDI method caji be applied i 'oad-band Doppler measurements.," The advantage of introducing an additional spectrograph into the system is that the visibility (or fringe contrast) is no longer limited by the bandpass as in the FTS case, which is the reason that the DFDI method can be applied in broad-band Doppler measurements."
+ ?/ discussed tie. possibility of an F'T5 workin[n] elu broad bau wy iutroduciug a low resolution post-cisperser and concluded that the FTS method is inferior (Nw a [actor between 1 aud 2) to DE method even aftereimiployiug au post-disperser., \citet{Mosser2003} discussed the possibility of an FTS working in broad band by introducing a low resolution post-disperser and concluded that the FTS method is inferior (by a factor between 1 and 2) to DE method even afteremploying an post-disperser.
+ This conclusio1 should be accepted with cautious because they compared an FTS with a post-cdisperser, This conclusion should be accepted with cautions because they compared an FTS with a post-disperser
+llere a is the standard. viscosity parameter for accretion disks. and h22ο is the verlical thickness of the disk with c; the sound speed.,"Here $\alpha$ is the standard viscosity parameter for accretion disks, and $h\approx c_s/\Omega$ is the vertical thickness of the disk with $c_s$ the sound speed."
+ The planets orbital eccentricity changes on the timescale where w is (he gaps width., The planet's orbital eccentricity changes on the timescale where $\w$ is the gap's width.
+ Equation(19).. adapted from Goldreich&Tremaine(1980).. is obtained by summiug contributions to «ερpiαἱ [rom first order resonances within a narrow ring separated from the planets position by an empty gap.," Equation, adapted from \cite{GOT80}, is obtained by summing contributions to $de_p/dt$ from first order resonances within a narrow ring separated from the planet's position by an empty gap."
+" Assuming lxeplerian rotation. thiev find that e, decavs."," Assuming Keplerian rotation, they find that $e_p$ decays."
+ However damping bv corotation resonances exceeds driving by Lindblad resonances by only a small. 4.676. margin.," However damping by corotation resonances exceeds driving by Lindblad resonances by only a small, $4.6\%$, margin."
+ The timescale quoted above for eccentricity change is that due to either corotation or Lindblad resonances acting separately., The timescale quoted above for eccentricity change is that due to either corotation or Lindblad resonances acting separately.
+" Conditions wider which this. rather than the net contributionfrom both twpes of resonance. is (he appropriate /. lo compare with /,;; are described in relsecisaturation.."," Conditions under which this, rather than the net contributionfrom both types of resonance, is the appropriate $t_e$ to compare with $t_{vis}$ are described in \\ref{sec:saturation}. ."
+" To elucidate the comparison between /, and /,;;. we relate w to a and AL,/AL. by balancing the viscous with the torque from the principal Lindblad resonances. to obtain Our typical parameters. a=107. h/r=0.04. M,/M,=10""7. give wosOn "," To elucidate the comparison between $t_e$ and $t_{vis}$, we relate $\w$ to $\alpha$ and $M_p/M_*$ by balancing the viscous with the torque from the principal Lindblad resonances, to obtain Our typical parameters, $\alpha=10^{-3}$, $h/r=0.04$, $M_p/M_*=10^{-3}$, give $\w\approx 0.4 r$."
+substituting this expression for w into equation vields For w/rz 1. eccentricity evolves faster (han semimajor axisprovided the disk is more massive (han the planet.," Substituting this expression for $\w$ into equation yields For $\w/r\approx 1$ , eccentricity evolves faster than semimajor axisprovided the disk is more massive than the planet."
+The forest is an important cosmological observable that probes matter density Ductuations in the IGM over a unique range of redshifts. scales and. environments.,"The forest is an important cosmological observable that probes matter density fluctuations in the IGM over a unique range of redshifts, scales and environments."
+ Many attempts have been mace to measure physical properties of the LGAL using. forest. data., Many attempts have been made to measure physical properties of the IGM using forest data.
+ The. two most common approaches are either. based on decomposing the information encoded in the transmitted [Dux via Voigt profile fitting or treating the ux as a continuous field with directly measurable statistical properties (e.g. 227773).," The two most common approaches are either based on decomposing the information encoded in the transmitted flux via Voigt profile fitting or treating the flux as a continuous field with directly measurable statistical properties (e.g. \citealt{rauch97,rauch98,theuns98,croft02,meiksin07}) )."
+ In the second approach. measurement of the zero. one. two-point or threc-point probability distribution functions (ic. the mean [ux level. the flux PDE. the Dux power and bispectrum) enable a variety of physical. properties to be explored.," In the second approach, measurement of the zero, one, two-point or three-point probability distribution functions (i.e. the mean flux level, the flux PDF, the flux power and bispectrum) enable a variety of physical properties to be explored."
+ The mean Hus level for example. is sensitive to the amplitude of the meta-galactic UV background (22)) while the Lux PDL is sensitive to the thermal evolution of the IGML (??)..," The mean flux level for example, is sensitive to the amplitude of the meta-galactic UV background \citealt{tytler04,bolt05}) ) while the flux PDF is sensitive to the thermal evolution of the IGM \citep{mcquinn09,bolton09}."
+ Phe tux power spectrum has been used to constrain cosmological parameters and the behaviour of dark matter at. small scales. (22?) and the Εαν bispectrum can be usec to search for signatures of non-gaussianities in the matter distribution (?)..," The flux power spectrum has been used to constrain cosmological parameters and the behaviour of dark matter at small scales \citep{vhs,seljak06,viel08} and the flux bispectrum can be used to search for signatures of non-gaussianities in the matter distribution \citep{viel09}."
+ Ideally. à given. IGM model described by a set of cosmological and astrophysical parameters should agree with all these statistics including the results from Voigt profile decomposition at the same time.," Ideally, a given IGM model described by a set of cosmological and astrophysical parameters should agree with all these statistics including the results from Voigt profile decomposition at the same time."
+ In practice. the interpretation of the data is complex. anc is heavily dependent on numerical simulations that incorporate the relevant physical ingredients. but have a limited: ἂνnamic range.," In practice, the interpretation of the data is complex and is heavily dependent on numerical simulations that incorporate the relevant physical ingredients, but have a limited dynamic range."
+ The data usec for these investigations consist mainly of two kinds of sets of QSO spectra: the SDSS low low signaltonoise sample. and. UVISS/VET or," The data used for these investigations consist mainly of two kinds of sets of QSO spectra: the SDSS low--resolution, low signal–to–noise sample and UVES/VLT or"
+2009).,.
+. In any other mocoel this represents a coincidence., In any other model this represents a coincidence.
+ Basically the same luminosity function behaviour is also seen in X-ray surveys (Bovleetal. 1993)., Basically the same luminosity function behaviour is also seen in X-ray surveys \citep{boyle93}.
+. Phere have been some reports that the X-ray LE is inconsistent with PLE (Hasingeretal.2001:Ciiacconial.2002:Alexanderct2003:Worsleyet2004). but these deviations are generally small and at low Ly and in a broad-brush way at least. PLE remains an approximate fit to the X-ray LE.," There have been some reports that the X-ray LF is inconsistent with PLE \citep{hasinger, giacconi,
+alexander, worsley} but these deviations are generally small and at low $L_X$ and in a broad-brush way at least, PLE remains an approximate fit to the X-ray LF."
+ PLE has always been an excellent phenomenological fit to the optical QSO LE (although see 3ongiornoetal.20072) but the question has been whether it has any physical meaning., PLE has always been an excellent phenomenological fit to the optical QSO LF (although see \citealt{bongiorno}) ) but the question has been whether it has any physical meaning.
+ Hf exactly true the PLE model would suggest that low z Sevfert Es would have as high a DII mass as high-z QSOs at fixed L/L because the BI fuelling rate at low recishift is small compared to that at high redshift (Marshall1985:Bovle.Shanks.&Peterson1988).," If exactly true the PLE model would suggest that low $z$ Seyfert I's would have as high a BH mass as high-z QSOs at fixed $L/L^*$ because the BH fuelling rate at low redshift is small compared to that at high redshift \citep{marshall, boyle88}."
+. A single DII mass independent. of QSO Duminosity. is not. however. a prediction of PLE at. fixed. redshift.," A single BH mass independent of QSO luminosity is not, however, a prediction of PLE at fixed redshift."
+ The equation above implies that for fixed 5 and « à QSO 10 brighter than L will produce a lO. bigger black hole mass of 107M. by z=0., The equation above implies that for fixed $\gamma$ and $\epsilon$ a QSO $10\times$ brighter than $L^*$ will produce a $10\times$ bigger black hole mass of $10^{10}{\rm M}_\odot$ by $z=0$.
+ Indeed. for à 10 QSO. from +=22 to z—1.5.ut of order the minimum QSO lifetime. a black-hole mass of !utLOPAL. will already have been created.," Indeed, for a $10L^*$ QSO, from $z=2.2$ to $z=1.5$, of order the minimum QSO lifetime, a black-hole mass of $5\times10^9{\rm
+M}_\odot$ will already have been created."
+ H£ we assume that the QSO has to be radiating at no more than Ecclington. this implies that the initial mass for such QSOs has to be LO°M..," If we assume that the QSO has to be radiating at no more than Eddington, this implies that the initial mass for such QSOs has to be $10^{9}{\rm M}_\odot$."
+ 1£ the fainter QSOs have a similar initial mass then this will minimise anv Mgg dillerence between bright and faint QSOs (see Fig. 3))., If the fainter QSOs have a similar initial mass then this will minimise any $M_{BH}$ difference between bright and faint QSOs (see Fig. \ref{fig:ple_gam_eps}) ).
+ OL course. if halo and BLL masses are completely decoupled. then there may. be no need. to accommodate the approximate luminosity independence of clustering in this way.," Of course, if halo and BH masses are completely decoupled, then there may be no need to accommodate the approximate luminosity independence of clustering in this way."
+ However. it could be that although halo/BLL mass ave clecouplec.tified halo ancl BIL masses could. still be correlated.," However, it could be that although halo/BH mass are decoupled, halo and BH masses could still be correlated."
+" I£ κο the Alew:Alnor, relation must evolve with redshift in the sense that the halo mass containing a particular Mgg must decrease as 2 increases."," If so, the $M_{BH}:M_{halo}$ relation must evolve with redshift in the sense that the halo mass containing a particular $M_{BH}$ must decrease as $z$ increases."
+ At.=22. all QSOs would then start with the same initial Alew and have the same halo mass to be consistent. with he clustering observations.," At $z=2.2$, all QSOs would then start with the same initial $M_{BH}$ and have the same halo mass to be consistent with the clustering observations."
+ Higher luminosity QSOs then accrete at a faster rate as in Fig., Higher luminosity QSOs then accrete at a faster rate as in Fig.
+ 3. but no luminosity. dependence of the QSO clustering is generated. because of he decoupling of halo and QSO/ealaxy merging rates anc he clustering of all QSOs grows at the long-lived rate., \ref{fig:ple_gam_eps} but no luminosity dependence of the QSO clustering is generated because of the decoupling of halo and QSO/galaxy merging rates and the clustering of all QSOs grows at the long-lived rate.
+ In the context of the original PLE model. QSOs at al uminosities at 2 could radiate at close to the IEddington imit and those at lower redshift then would radiate a increasingly sub-ISddington rates. up to z30..," In the context of the original PLE model, QSOs at all luminosities at $z\approx2$ could radiate at close to the Eddington limit and those at lower redshift then would radiate at increasingly sub-Eddington rates, up to $\approx30\times$."
+ In the above PLE model case where all the initial Mgg must now be he same to fit the clusteringresults. if we assume that a 2& the brightest QSO is raciating at Edclineton. then at the same redshift the fainter QSOs have now to radiate at =10 lower than the Edcington rate.," In the above PLE model case where all the initial $M_{BH}$ must now be the same to fit the clusteringresults, if we assume that at $z\approx2$ the brightest QSO is radiating at Eddington, then at the same redshift the fainter QSOs have now to radiate at $\approx10\times$ lower than the Eddington rate."
+ At 2&0. these intrinsically faint QSOs have thus now to be radiating at zz300. below Eddington.," At $z\approx0$, these intrinsically faint QSOs have thus now to be radiating at $\approx300\times$ below Eddington."
+ This may be physically feasible in that standard. optically thick. physically thin. accretion discs can accrete at down to z14 of the Exldington rate (€. Done. priv.," This may be physically feasible in that standard optically thick, physically thin, accretion discs can accrete at down to $\approx1$ of the Eddington rate (C. Done, priv."
+ conum.), comm.).
+ 3clow this rate the accretion disc is unstable to evaporation., Below this rate the accretion disc is unstable to evaporation.
+ Thus the -27 mag QSOs at high redshift and the -22 mag Sevlert Vs locally can both be powered. by thin disc accretion., Thus the -27 mag QSOs at high redshift and the -22 mag Seyfert I's locally can both be powered by thin disc accretion.
+ Thus there appears to be a reasonable range of QSO luminosities where the PLE or single-DII mass moclels can operate within the context of a thin clise accretion moclel., Thus there appears to be a reasonable range of QSO luminosities where the PLE or single-BH mass models can operate within the context of a thin disc accretion model.
+ Vhe mean BIE mass measured in low luminosity SDSS AGN by Lleckmanetal.(2004) via the Alewτσ relation is zLOM. and this seems in less good. agreement with the =10ΛΙ. PLE prediction than the results of Laor (1990)..," The mean BH mass measured in low luminosity SDSS AGN by \citet{heckman}
+ via the $M_{BH}:\sigma$ relation is $\approx10^8 M_\odot$ and this seems in less good agreement with the $\approx10^9 M_\odot$ PLE prediction than the results of \cite{laor}. ."
+ Llowever. as we have already noted. the SDSS AGN comprise some of the galaxy population whereas," However, as we have already noted, the SDSS AGN comprise some of the galaxy population whereas"
+constrain “lp and z.,constrain $A_{B}$ and $\varepsilon$.
+ For this example. we chose constraints consistent with Ap=0.2 and 5=0.3.," For this example, we chose constraints consistent with $A_{B}=0.2$ and $\varepsilon=0.3$."
+" The solid line is a locus of constant dj: the dotted line is the locus of coustaut {τω the dashed line is a locus of constant 2,,/7y.", The solid line is a locus of constant $A_{B}$; the dotted line is the locus of constant $T_{d}/T_{0}$; the dashed line is a locus of constant $T_{n}/T_{0}$.
+ From this figure it is clear that the measurements complement cach otler: nieasurng two of the three quantities (Boud albedo. effective dav-side or uight-side temperatures) uniquely determines the planets albedo aud circulation efficiency.," From this figure it is clear that the measurements complement each other: measuring two of the three quantities (Bond albedo, effective day-side or night-side temperatures) uniquely determines the planet's albedo and circulation efficiency."
+ When observations lave some associated uncertainty. they define a swath through the Ap 5 plane.," When observations have some associated uncertainty, they define a swath through the $A_{B}$ $\varepsilon$ plane."
+ We οσο. by cousidering all the photometric observations of short-period exoplanets published through November 2010. summarized in Table 1.," We begin by considering all the photometric observations of short-period exoplanets published through November 2010, summarized in Table \ref{observations}."
+ We have discarded photometric observations of non-transiting plauets because of their πιο radius aud orbitalinclination’., We have discarded photometric observations of non-transiting planets because of their unknown radius and orbital.
+.. This leaves us with 21 transiting exoplauets for which there are observations iu atf least one waveband at superior conjunction. and im some cases in multiple wavebauds aud at multiple planetary phases.," This leaves us with 24 transiting exoplanets for which there are observations in at least one waveband at superior conjunction, and in some cases in multiple wavebands and at multiple planetary phases."
+ Stellar aud planetary data are taken from the Exoplanet Encyclopedia (exoplanct.cu). aud references herein.," Stellar and planetary data are taken from the Exoplanet Encyclopedia (exoplanet.eu), and references therein."
+ We repeated parts of the analvsis with the Exoplanet Data Explorer database (exoplanets.org) aud ound identical results. within the uucertaiuties.," We repeated parts of the analysis with the Exoplanet Data Explorer database (exoplanets.org) and found identical results, within the uncertainties."
+ When he stellar data are not available. we have assumed. vpical parameters for the appropriate spectral class. and solar metallicity.," When the stellar data are not available, we have assumed typical parameters for the appropriate spectral class, and solar metallicity."
+ Insofar as we are ouly concerned with the broadband brightucsses of the stars. our results should not depend scusitively on the iuput stellar parameters.," Insofar as we are only concerned with the broadband brightnesses of the stars, our results should not depend sensitively on the input stellar parameters."
+ Knowing the stars Tig. logg and |Fe/II]. we use the PIIOENIN/NostGeu stellay spectrum erids (lauschildtetal.1999). to determine their briehtuess teniperatures at the observed bandpasses.," Knowing the stars' $T_{\rm eff}$, $\log g$ and [Fe/H], we use the PHOENIX/NextGen stellar spectrum grids \citep{Hauschildt_1999} to determine their brightness temperatures at the observed bandpasses."
+ At cach wavebaud for which eclipse or phase observatious have been obtained. we determine the ratio of the stellar flux to the blackbody Hux at that exid stars Tig.," At each waveband for which eclipse or phase observations have been obtained, we determine the ratio of the stellar flux to the blackbody flux at that grid star's $T_{\rm eff}$."
+ We then apply this factor to he Tige of the observed star., We then apply this factor to the $T_{\rm eff}$ of the observed star.
+ Tt is worth noting that the choice of stellar model leads ο systematic uncertainties in the planctary brightucss hat are of order the photometric unucertaimties., It is worth noting that the choice of stellar model leads to systematic uncertainties in the planetary brightness that are of order the photometric uncertainties.
+ For exaluple. Cliistiansenctal.(2010) use stellar inodels or ILAT-P-7 from I&urucz(2005).. while we use those of Uauschildtetal.(1999).," For example, \cite{Christiansen_2009} use stellar models for HAT-P-7 from \cite{Kurucz_2005}, while we use those of \cite{Hauschildt_1999}."
+. The resulting 8 jan brightuess eniperatures for ILAT-P-7b differ by as much as 600 Ix. or slightly wore than Lo.," The resulting 8 $\mu$ m brightness temperatures for HAT-P-7b differ by as much as 600 K, or slightly more than $1\sigma$."
+ Our uniform use of Hauschildtetal.(1999) models should alleviate this problem. 10Wever.," Our uniform use of \cite{Hauschildt_1999} models should alleviate this problem, however."
+" The plauct’s albedo aud recirculation cficicney govern its effective day-side aud might-side temperatures. Ty aud T4. respectively,"," The planet's albedo and recirculation efficiency govern its effective day-side and night-side temperatures, $T_{\rm d}$ and $T_{\rm n}$, respectively."
+ Observationallv. we can ouly measure he brightness temperature. ideally at a umuber of different wavelengths: Πλ).," Observationally, we can only measure the brightness temperature, ideally at a number of different wavelengths: $T_{\rm b}(\lambda)$."
+ Tf one. knew.priori. he cimergent spectrum of a plauet. oue could trivially convert a single brightuess temperature to an effective eniperature.," If one knew, the emergent spectrum of a planet, one could trivially convert a single brightness temperature to an effective temperature."
+ Alternatively. if observations were obtained ata number of waveleneths bracketing the planuct’s dackbody peal. it would be possible to estimate he planet's bolometric flux and hence its effective cluperature in a modceLindependeut way (c.g..Barua2008).," Alternatively, if observations were obtained ata number of wavelengths bracketing the planet's blackbody peak, it would be possible to estimate the planet's bolometric flux and hence its effective temperature in a model-independent way \citep[e.g.,][]{Barman_2008}."
+. We adopt the latter empirical approach of converting observed flux ratios iuto. brightness temiperatures. hen using these to estimate the planet's effective eniperature.," We adopt the latter empirical approach of converting observed flux ratios into brightness temperatures, then using these to estimate the planet's effective temperature."
+ The secondary eclipse depth in some waveband divided by the transit depth is a direct neasure of the ratio of the plauct’s day-side intensity to he stars intensity at that wavelength. ofA}.," The secondary eclipse depth in some waveband divided by the transit depth is a direct measure of the ratio of the planet's day-side intensity to the star's intensity at that wavelength, $\psi(\lambda)$."
+ Kuowiug he stars brightuess temperature at a eiven wavceleusth. it ds possible to compute the apparent brightuess eniperature of the plauct’s day side: Ou the Ravleigh-Jeaus tail. the fractional uncertaiuty in the brightuess temperature is roughly equal to the ractional uncertainty in the eclipse depth: ou the Wien sd. the fractional error on brightuess teniperature can be smaller because the flux is very sensitive to cluperature.," Knowing the star's brightness temperature at a given wavelength, it is possible to compute the apparent brightness temperature of the planet's day side: On the Rayleigh-Jeans tail, the fractional uncertainty in the brightness temperature is roughly equal to the fractional uncertainty in the eclipse depth; on the Wien tail, the fractional error on brightness temperature can be smaller because the flux is very sensitive to temperature."
+ By the same token. a secondary eclipse depth aud phase variation amplitude at a given waveleneth can be conibined to obtain a measure of the planets might-sicde xiehtuess temperature at that svavebaud.," By the same token, a secondary eclipse depth and phase variation amplitude at a given wavelength can be combined to obtain a measure of the planet's night-side brightness temperature at that waveband."
+ Since the albedo and recirculation eficicney of the alanet are not mown ahead of tine. it is not immediately obvious which wavelengths are scusitive to reflected ight and which are donmüuated by thermal cussion.," Since the albedo and recirculation efficiency of the planet are not known ahead of time, it is not immediately obvious which wavelengths are sensitive to reflected light and which are dominated by thermal emission."
+ For each planet. we compute the expected blackbody »eak if the planet has no albedo aud no recirculation of energv. Azy=28ON/T. µια. Dusofar as real plauets will have non-zero albedo and non-zero recirculation. he dav side should never reach 7.y. aud the actual spectral energy distribution will peak at καστ] longer waveleugths.," For each planet, we compute the expected blackbody peak if the planet has no albedo and no recirculation of energy, $\lambda_{\varepsilon=0} = 2898/T_{\varepsilon=0}$ $\mu$ m. Insofar as real planets will have non-zero albedo and non-zero recirculation, the day side should never reach $T_{\varepsilon=0}$, and the actual spectral energy distribution will peak at slightly longer wavelengths."
+ The coolest planet in our sample. C1 136 would exhibit a blackbody peak at A.9=3.1 jan. while the hottest plauct we cousider. WASP-12b. has Αν.y=0.9 gon. In practice this meaus that eround-based near-IR and space-based mid-IR (e.g. Spitcer)) observations are assumed to measure thermal euission. while space-based optical observations (MOST. CoRoT. Isepler) may be coutamunated by reflected starlight.," The coolest planet in our sample, Gl 436b, would exhibit a blackbody peak at $\lambda_{\varepsilon=0}=3.1$ $\mu$ m, while the hottest planet we consider, WASP-12b, has $\lambda_{\varepsilon=0}=0.9$ $\mu$ m. In practice this means that ground-based near-IR and space-based mid-IR (e.g., ) observations are assumed to measure thermal emission, while space-based optical observations (MOST, CoRoT, Kepler) may be contaminated by reflected starlight."
+ In Figure 2. we demonstrate two alternative echuiques το. couvert oan array of brightness eniperatures. 2(A}. iuto an estimate of a planets effective temperature. Digg.," In Figure \ref{tres3_2pi_TiO.flx}, , we demonstrate two alternative techniques to convert an array of brightness temperatures, $T_{\rm b}(\lambda)$, into an estimate of a planet's effective temperature, $T_{\rm eff}$."
+ The solid black line shows a nodel spectrum of thermal cimissiou from Fortueyetal. (2008).. with an effective temperature of Tig—1911 Is shown with the black dashed Lue.," The solid black line shows a model spectrum of thermal emission from \cite{Fortney_2008}, with an effective temperature of $T_{\rm eff} = 1941$ K shown with the black dashed line."
+ The expected dlackbodxy peak of the planet is marked with a vertical dotted line., The expected blackbody peak of the planet is marked with a vertical dotted line.
+" The red points are the expected brightuess eniperatures in the J. IL. aud EK, bands(crosses). as well as the IRAC (asterisks) and MIPS. (diaunnoud) instruments ou (Fazioctal.2001:Riekeet 2001).."," The red points are the expected brightness temperatures in the J, H, and $_{\rm s}$ bands(crosses), as well as the IRAC (asterisks) and MIPS (diamond) instruments on \citep{Fazio_2004, Rieke_2004, Werner_2004}. ."
+ Since the majority of the, Since the majority of the
+and the cussionweighted metallicity. in the |l]2] keV band.,"and the emission–weighted metallicity, in the [1–2] keV band."
+ Then. we compute the corresponding MERAL spectrmu for this teuperatiure and metallicity. to be cored with the actual svuthetic spectrum in the [0.21) keV band (ye use in the following the solar photospheric abtvdance by Anders Crevesse 1989. when expressing mctallici vdn solar units).," Then, we compute the corresponding MEKAL spectrum for this temperature and metallicity, to be compared with the actual synthetic spectrum in the [0.2–1] keV band (we use in the following the solar photospheric abundance by Anders Grevesse 1989, when expressing metallicity in solar units)."
+ Ccarly. our approach is correc as loug as enussonweighted measures coincide with the corresponding quautities obtained from spectral fitting.," Clearly, our approach is correct as long as emission–weighted measures coincide with the corresponding quantities obtained from spectral fitting."
+ MNazzotta ct al. (, Mazzotta et al. (
+"2001) have recently pointed out that the ο1uplex ICAL teiiperature structure m simulated chisters c‘auses the spectralfitting temperatures to be about 20 yer cent lower thai T,Ww (see also. Alathiesen Evrard 2001).",2004) have recently pointed out that the complex ICM temperature structure in simulated clusters causes the spectral–fitting temperatures to be about 20 per cent lower than $T_{\rm ew}$ (see also Mathiesen Evrard 2001).
+ They provide au expression for a specroscopiclike temperature. which represents an effective recipe to compte the actual spectral temperature 1 simmtlations.," They provide an expression for a spectroscopic–like temperature, which represents an effective recipe to compute the actual spectral temperature in simulations."
+ However. the fitting πιοοι. even by Mazzotta ct al.," However, the fitting function given by Mazzotta et al."
+ ds xpecific to the response function and eunergev coverage o| the Chandra and NATALNewton detectors. therefore iof. applicable o the ROSATPSPC.," is specific to the response function and energy coverage of the Chandra and XMM–Newton detectors, therefore not applicable to the ROSAT–PSPC."
+ Furthermore. no such effective recipe las been calibrated πο far to estimate the spectroscopic metallicity in simulated chsters.," Furthermore, no such effective recipe has been calibrated so far to estimate the spectroscopic metallicity in simulated clusters."
+ Finally. he spectroscopiclike temperature is well «efined ouly for clusters hotter than 3 keV. while a significant fraction of he simulated clusters in our set has lower temperature (see Table 1).," Finally, the spectroscopic–like temperature is well defined only for clusters hotter than 3 keV, while a significant fraction of the simulated clusters in our set has lower temperature (see Table 1)."
+ For these reasons. we prefer to use here eqs.(2)) and (3)) for the definition of teperature axd metallicity.," For these reasons, we prefer to use here \ref{eq:tew}) ) and \ref{eq:zew}) ) for the definition of temperature and metallicity."
+" We defer to a forthcoming paper a nkπο detailed analysis, based on svuthetic NAALNewton n aud MOS spect of Sinn clisters, to 1ivestie:ite the detectability of a soft excesslated muder reaistic observi18o conditions (Cheng et al."," We defer to a forthcoming paper a more detailed analysis, based on synthetic XMM-Newton pn and MOS spectra of simulated clusters, to investigate the detectability of a soft excess under realistic observing conditions (Cheng et al.,"
+ in preparation)., in preparation).
+ The values of Tia ane Ζω in the |]2| τον baAC are given iu Table 1 for all our clusters., The values of $T_{\rm ew}$ and $Z_{\rm ew}$ in the [1–2] keV band are given in Table 1 for all our clusters.
+ The resulti18o metallicity values are sinalor than the typical observe Fe abundance. ~0.3Z.. (6... Arnaud et al.," The resulting metallicity values are smaller than the typical observed Fe abundance, $\sim 0.3
+Z_{\odot}$ (e.g., Arnaud et al."
+ 2007: Usehe al., 2001; Ikebe et al.
+ 2002: Damugartuer et aο, 2002; Baumgartner et al.
+ 2003)., 2003).
+ A careful study of t ICM inetalliciv would require includiue in simulations t coutribuion of the SN-Ia aud accounutiine for the cdiffereu vields different elements (e.g... Toruatore et al.," A careful study of the ICM metallicity would require including in simulations the contribution of the SN-Ia and accounting for the different yields of different elements (e.g., Tornatore et al."
+ 200, 2004).
+" Tha prie. tils lua Leprsent a ndation of the preseu analysis. since ποτά] iesre expectec ο give a significa coutributiou to the toal emissivitv iu he soft part oft1C spectra where we are se‘sine for the excess,"," In principle, this may represent a limitation of the present analysis, since metal lines are expected to give a significant contribution to the total emissivity in the soft part of the spectrum, where we are seeking for the excess."
+ In order to test whether our final resuls are affect(0 by the uncertaiu descrition of the ICN metal euricliuieit. we have verified by how much our results ¢change when we assume cit101) Z Yor Z0257. oei the MERAL spectymu of a plasma with temperature JT=2 keV. We find that he contribu]ion to the total οwussion from metal lines turus out ohe~3 per cent iu he |0.21] keV band. this tiny difference beire due to the lack of prominent lines in he above energv range.," In order to test whether our final results are affected by the uncertain description of the ICM metal enrichment, we have verified by how much our results change when we assume either $Z=0$ or $Z=0.25 Z_{\odot}$ in the MEKAL spectrum of a plasma with temperature $T=2$ keV. We find that the contribution to the total emission from metal lines turns out to be $\sim 3$ per cent in the [0.2–1] keV band, this tiny difference being due to the lack of prominent lines in the above energy range."
+ Therefore. we expect that our fial results shoul be lavecly insensitive to the approximate treatiueut of the ICAL meted enrichient.," Therefore, we expect that our final results should be largely insensitive to the approximate treatment of the ICM metal enrichment."
+for ΠοZ:10°. aud Ty can become the dominant hydrogen species at the plotosphere.,"for $\rm He/H\wig>10^{3}$, and $\rm H_2$ can become the dominant hydrogen species at the photosphere."
+" The most important sources of opacity in helimmaaich white cawarfatmospheres are ο free-free, Ravleigh scattering. aud H5Ue CIA (ILuseu2001)."," The most important sources of opacity in helium-rich white dwarfatmospheres are $\rm He^{-}$ free-free, Rayleigh scattering, and $\rm H_{2}-He$ CIA \citep{HN}."
+. Several physical and chemical effects that alter these ορα sources have already been discussed: the change in the free-free opacity aud the Ravleigh scattering of helium. as caused by the stroug correlations iu the deuse fluid (elesiasetal.2002).. the change iu the nuunmber deusitv of free electrons (Ixowalskictal.2005:Berecronetal. 1995).. the presence of heavy eleineuts (Berecron 2001).. and the formation of trace specics like Te (Maloetal.1999). and Uell! (Παςetal.2001)..," Several physical and chemical effects that alter these opacity sources have already been discussed: the change in the free-free opacity and the Rayleigh scattering of helium, as caused by the strong correlations in the dense fluid \citep{IRS}, the change in the number density of free electrons \citep{KSM05,BSW}, the presence of heavy elements \citep{Bergeron01}, , and the formation of trace species like $\rm He_{2}^{+}$ \citep{ML} and $\rm HeH^{+}$ \citep{HR}."
+ The effect ou ποιασα] chemistry on the IT» dissociation on the opacity is shown in the right haud panels of Figure |., The effect on non-ideal chemistry on the $\rm H_2$ dissociation on the opacity is shown in the right hand panels of Figure 4.
+ The increase m umber deusity of molecular hivdrogeu results iu an increase of the ΠοTe CTA opacity., The increase in number density of molecular hydrogen results in an increase of the $\rm H_2-He$ CIA opacity.
+ This follows from the linear depeudence of IT;Ie CTA opacity on ay., This follows from the linear dependence of $\rm H_2-He$ CIA opacity on $n_{\rm H_2}$.
+ Ou the other laud. the II bouud-free and frec-free opacities are reduced. due to the decrease in the abuudance of atomic hydrogen.," On the other hand, the $\rm H^-$ bound-free and free-free opacities are reduced, due to the decrease in the abundance of atomic hydrogen."
+" Based on the atmospheric structures used here. the effects of the nou-deal chemistry of Ilo iu dense helmnu are maximal for IHe/I ~100,"," Based on the atmospheric structures used here, the effects of the non-ideal chemistry of $_2$ in dense helium are maximal for He/H $\sim 10^3$."
+ This arises from a competition between the need for a high UWe/T ratio to increase the uon-ideal effects by inereasing the density at the photosphere (Fig., This arises from a competition between the need for a high He/H ratio to increase the non-ideal effects by increasing the density at the photosphere (Fig.
+ 2b) and the need for a high enough hydrogen coutent in the atmosphere so that II» (or II) contributes to the total opacity (Fie., 2b) and the need for a high enough hydrogen content in the atmosphere so that $_2$ (or H) contributes to the total opacity (Fig.
+ Ll. night haud sido).," 4, right hand side)."
+ The impact of the non-ideal dissociation equilibrium of II on the svuthetic spectrum of a TWe-rich white wart model is shown in Figure 5., The impact of the non-ideal dissociation equilibrium of $\rm H_{2}$ on the synthetic spectrum of a He-rich white dwarf model is shown in Figure 5.
+ For a white dwarf atinosphere model of Tig=100019. Ίουy=8. aud ΠοΠ=10°. the vise in abundance of II» increases significantly the opacity iu the infrared. causing a redistribution of the fiux toward shorter wavelengths.," For a white dwarf atmosphere model of $T_{\rm eff}\rm=4000K$, $\rm log \it \ g\rm=8$, and $\rm He/H=10^{3}$, the rise in abundance of $\rm H_2$ increases significantly the opacity in the infrared, causing a redistribution of the flux toward shorter wavelengths."
+ This effect on the spectrum of cool white dwarts is largest for Tig=1000 15001. where IT; is partially dissociated.," This effect on the spectrum of cool white dwarfs is largest for $T_{\rm eff}\rm =4000K-4500K$ , where $\rm H_2$ is partially dissociated."
+ At lower effective. temperatures hwvdrogen exists mostly iu inolecular form aud the effect of the uon-idoeal dissociation equilibrium ou the spectrua vanishes at Tig~ 30001., At lower effective temperatures hydrogen exists mostly in molecular form and the effect of the non-ideal dissociation equilibrium on the spectrum vanishes at $T_{\rm eff}\rm \sim 3000K$ .
+ The determination of the ΠοἩ composition of very cool Ποιο white dwarfs depeucds mostly on the relative IeIl;CIA. aud Ile frec-free opacities (Fig.," The determination of the $\rm He/H$ composition of very cool He-rich white dwarfs depends mostly on the relative $\rm He-H_2$CIA, and $\rm He^-$ free-free opacities (Fig."
+ Ll. melthaud side). Since the," 4, righthand side), Since the"
+detector have been evaluated by using the Moon shadow. Le. the deficit of cosmic rays in the Moon direction.,"detector have been evaluated by using the Moon shadow, i.e. the deficit of cosmic rays in the Moon direction."
+ The shape of the shadow provides a measurement of the detector Point Spread Function (PSF). and its position allows the individuation of possible pointing biases.," The shape of the shadow provides a measurement of the detector Point Spread Function (PSF), and its position allows the individuation of possible pointing biases."
+" ARGO-YBJ observes the Moon shadow with a sensitivity of about 10 standard deviations per month for events with a multiplicity Noo 240 and zenith angle 7«503. corresponding to a proton median energy E, ~1.8 TeV (DiSciascioetal.2008)."," ARGO-YBJ observes the Moon shadow with a sensitivity of about 10 standard deviations per month for events with a multiplicity $_{pad} \geq$ 40 and zenith angle $\theta <$ $^{\circ}$, corresponding to a proton median energy $_p\sim$ 1.8 TeV \cite{DiS08}."
+. According to the Moon shadow data. the PSF of the detector is Gaussian for Να 22100. while for lower multiplicities it can. be deseribed with an additional Gaussian. which contributes for about20%.," According to the Moon shadow data, the PSF of the detector is Gaussian for $_{pad} \ge$ 100, while for lower multiplicities it can be described with an additional Gaussian, which contributes for about."
+. When the PSF is a Gaussian with r.m.s. 7.," When the PSF is a Gaussian with r.m.s. $\sigma$,"
+ the opening angle c containing ~71.5% of the events maximizes the signal to background ratio for a point source with a uniform background. and is equal to 1.58 c.," the opening angle $\psi$ containing $\sim$ $\%$ of the events maximizes the signal to background ratio for a point source with a uniform background, and is equal to 1.58 $\sigma$."
+" The semi-aperture c is found to be 2.59""3- 0.167. cx 0.147 and 1.04°+ 0.12"" for 240. 100 and 300. respectively. in agreement with expectationsN,4,; from Monte Carlo simulations."," The semi-aperture $\psi$ is found to be $^{\circ} \pm$ $^{\circ}$, $^{\circ} \pm$ $^{\circ}$ and $^{\circ} \pm$ $^{\circ}$ for $_{pad}\geq$ 40, 100 and 300, respectively, in agreement with expectations from Monte Carlo simulations."
+ This measured angular resolution refers to cosmic ray- air showers., This measured angular resolution refers to cosmic ray-induced air showers.
+" The angular resolution for ~-inducec events has been evaluated by simulations and results smaller by ~30-40%. depending on N,,. due to the better defined time profile of the showers."," The angular resolution for $\gamma$ -induced events has been evaluated by simulations and results smaller by $\sim$ $\%$, depending on $_{pad}$, due to the better defined time profile of the showers."
+ The relation between the observed pad/strip multiplicity spectrum and the primary energy spectrum has been studiec with cosmic ray showers. by means of a full Monte Carlo simulation. including the CORSIKA code (Hecketal.1998) to describe the shower development in the atmosphere. anc a code based on the GEANT package (GEANT1993). to simulate the detector response.," The relation between the observed pad/strip multiplicity spectrum and the primary energy spectrum has been studied with cosmic ray showers, by means of a full Monte Carlo simulation, including the CORSIKA code \cite{Hec98} to describe the shower development in the atmosphere, and a code based on the GEANT package \cite{Gea93} to simulate the detector response."
+ Primary particles have beer sampled from the energy spectra of Hórrandel (2003)., Primary particles have been sampled from the energy spectra of Hörrandel (2003).
+ The measured strip multiplicity spectrum is 1n. good agreement with the one predicted by the simulation (Aiellietal.2009b)., The measured strip multiplicity spectrum is in good agreement with the one predicted by the simulation \cite{Aie09b}.
+". The dataset for the analysis of Mrk421 presented in this paper contains all showers with N,,4 40 and zenith angle less than 40.", The dataset for the analysis of Mrk421 presented in this paper contains all showers with $_{pad} \geq$ 40 and zenith angle less than $^{\circ}$.
+ No event selection based on the shower core position and no gamma-hadron discrimination have been applied in this work., No event selection based on the shower core position and no gamma-hadron discrimination have been applied in this work.
+ A sky map in celestial coordinates (right ascension and declination) with 0.17?«0.17 bin size. centered on the source location. is filled with the detected events.," A sky map in celestial coordinates (right ascension and declination) with $^{\circ}\times$ $^{\circ}$ bin size, centered on the source location, is filled with the detected events."
+ In order to extract the excess of 7 rays coming from the source. the cosmic ray background must be carefully estimated and subtracted from the event map.," In order to extract the excess of $\gamma$ rays coming from the source, the cosmic ray background must be carefully estimated and subtracted from the event map."
+ The background is evaluated with the nethod (Alexandreasetal.1993)., The background is evaluated with the method \citep{Ale92}.
+". For each detected event. ""fake"" events are generated by replacing the original arrival time with new ones. randomly selected from a buffer that spans a time T of data taking."," For each detected event, N ""fake"" events are generated by replacing the original arrival time with new ones, randomly selected from a buffer that spans a time T of data taking."
+ We chose T ~ 3 hours to ninimize the systematic effects due to the environmental parameters variations., We chose T $\sim$ 3 hours to minimize the systematic effects due to the environmental parameters variations.
+ Changing the time. the fake events naintain the same declination of the original event. but have a different right ascension.," Changing the time, the fake events maintain the same declination of the original event, but have a different right ascension."
+ With these events a new sky nap (background map) is built., With these events a new sky map (background map) is built.
+ The number of fake events generated for each event is N = «15 «cos(ó2c. where cis the radius of the observationalwindow in degrees (see below) and 9 is the declination of the source.," The number of fake events generated for each event is N = $\times$ $\times$ $\delta$ $\psi$, where $\psi$ is the radius of the observationalwindow in degrees (see below) and $\delta$ is the declination of the source."
+ In this way the average number of fake events falling in the observational window Is 1., In this way the average number of fake events falling in the observational window is $\sim$ 1.
+" The two maps are then ""integrated"" over a circular area of radius c. Le. every bin is filled with the content of all bins whose center has an angular distance less than «c from its center. with c — 1.77. 0.9° and 0.6° for N44240. 100 and 300. respectively."," The two maps are then ""integrated"" over a circular area of radius $\psi$, i.e. every bin is filled with the content of all bins whose center has an angular distance less than $\psi$ from its center, with $\psi$ = $^{\circ}$, $^{\circ}$ and $^{\circ}$ for $_{pad}\geq$ 40, 100 and 300, respectively."
+" Finally the integrated background map is subtracted to the corresponding integrated event map. obtaining the ""source map"". where for every bin the statistical significance of the excess Is calculated."," Finally the integrated background map is subtracted to the corresponding integrated event map, obtaining the ""source map"", where for every bin the statistical significance of the excess is calculated."
+ With this procedure. however. since also the source events are used in the time swapping procedure. the obtained background at the source position is slightly overestimated. and the signal underestimated.," With this procedure, however, since also the source events are used in the time swapping procedure, the obtained background at the source position is slightly overestimated, and the signal underestimated."
+" This underestimation increases with the observational window size. ranging from ~4 to 10% of the signal. depending on the Nj, interval."," This underestimation increases with the observational window size, ranging from $\sim$ 4 to $\%$ of the signal, depending on the $_{pad}$ interval."
+ The observed event rate is then corrected using the appropriate factor., The observed event rate is then corrected using the appropriate factor.
+ The whole procedure has been tested with the Crab Nebula. the standard candle for VHE astronomy.," The whole procedure has been tested with the Crab Nebula, the standard candle for VHE astronomy."
+" At the Yangbajing latitude the Crab culminates at a zenith angle (yi,28.17 and it is observableevery day for 5.8 hours with a zenith angle 0«AQ.", At the Yangbajing latitude the Crab culminates at a zenith angle $\theta_{culm} = 8.1^{\circ}$ and it is observableevery day for 5.8 hours with a zenith angle $\theta <$ $^{\circ}$.
+" The Crab Nebula was observed from 2007 December 13 to 2009 August 8. for a total of 3150 on-source hours. obtaining a signal with à statistical significance of 7.6 standard deviations for N,.¢240."," The Crab Nebula was observed from 2007 December 13 to 2009 August 8, for a total of 3150 on-source hours, obtaining a signal with a statistical significance of 7.6 standard deviations for $_{pad}\geq$ 40."
+ The average number of gamma rays detected per day is 156.6220.6 for ρω240., The average number of gamma rays detected per day is $\pm$ 20.6 for $_{pad}\geq$ 40.
+ To evaluate the energy spectrum. we simulate a source in the sky following the daily path of the Crab Nebula. and estimate the number of events expected in different ρω intervals. as a function of the spectrum parameters.," To evaluate the energy spectrum, we simulate a source in the sky following the daily path of the Crab Nebula, and estimate the number of events expected in different $_{pad}$ intervals, as a function of the spectrum parameters."
+ Assuming a power law spectrum in the 0.1-80 TeV energy range. the best fit to the data is dN/dE = (4160.6) « 107! (E/I Τον)Το photons ems! TeV!. in agreement with our previous measurement (Vernettoetal.2009) and with observations by other detectors. such as H.E.S.S. (Aharonianetal.2006).. MAGIC (Albertetal.2008). and Tibet AS-~ (Amenomorietal.2009).," Assuming a power law spectrum in the 0.1-80 TeV energy range, the best fit to the data is dN/dE = $\pm$ 0.6) $\times$ $^{-11}$ (E/1 $^{-2.7\pm 0.2}$ photons $^{-2}$ $^{-1}$ $^{-1}$, in agreement with our previous measurement \citep{Ver09} and with observations by other detectors, such as H.E.S.S. \citep{Aha06}, MAGIC \citep{Alb08} and Tibet $\gamma$ \citep{Ame09}."
+. This result confirms the reliability of the simulation procedure and of the energy calibration of the detector., This result confirms the reliability of the simulation procedure and of the energy calibration of the detector.
+ Concerning the energy range sampled in the Crab Nebula measurement. about 84% of the detected events comes from primary photons of energies greater than 300 GeV. while only 8% comes from primaries above 10 TeV. The same analysis was performed for Mrk421.," Concerning the energy range sampled in the Crab Nebula measurement, about $\%$ of the detected events comes from primary photons of energies greater than 300 GeV, while only $\%$ comes from primaries above 10 TeV. The same analysis was performed for Mrk421."
+ This source culminates at the ARGO-YBJ location at a zenith angle Üig=8.17. and it is observable every day for 6.4 hours with a zenith angle 7< 40°.," This source culminates at the ARGO-YBJ location at a zenith angle $\theta_{culm} = 8.1^{\circ}$, and it is observable every day for 6.4 hours with a zenith angle $\theta <$ $^{\circ}$."
+ We evaluate the Mrk421 spectrum from day 41 to 180 of 2008. when the X-ray flux showed the nost intense flares.," We evaluate the Mrk421 spectrum from day 41 to 180 of 2008, when the X-ray flux showed the most intense flares."
+ In this period (754 observation hours) the signal has a statistical significance of 5.8 standard deviations., In this period (754 observation hours) the signal has a statistical significance of 5.8 standard deviations.
+ We assume a power law spectrun multiplied by an exponential factor e7*'* to take into account the absorption of gamma rays in the Extragalactic Background Light (EBL). with the values of the optical depth 7(E) given by Raue & Mazin (2008).," We assume a power law spectrum multiplied by an exponential factor $^{-\tau(E)}$ to take into account the absorption of gamma rays in the Extragalactic Background Light (EBL), with the values of the optical depth $\tau(E)$ given by Raue $\&$ Mazin (2008)."
+" The best fit spectrum obtained is: dN/dE = (3.00.5) « 107 (E/1.5 Τον44"" e-7*! photons em s! Τον].", The best fit spectrum obtained is: dN/dE = $\pm0.5$ ) $\times$ $^{-11}$ (E/1.5 $^{-2.4\pm0.3}$ $^{-\tau(E)}$ photons $^{-2}$ $^{-1}$ $^{-1}$.
+" The integral flux above | TeV is (3.6£0.6) « 107""! photons em s. almost twice the Crab Nebula one. ie. 2.1 1011 photons em s!. according to Aharonian et al. ("," The integral flux above 1 TeV is $\pm0.6$ ) $\times$ $^{-11}$ photons $^{-2}$ $^{-1}$ , almost twice the Crab Nebula one, i.e. 2.1 $\times$ $^{-11}$ photons $^{-2}$ $^{-1}$ , according to Aharonian et al. ("
+2006).,2006).
+ The values of the spectral index and of the gamma ray flux averaged over this 140 days period. support the correlation between spectral hardness and flux intensity reported. by Krennrich et al. (," The values of the spectral index and of the gamma ray flux averaged over this 140 days period, support the correlation between spectral hardness and flux intensity reported by Krennrich et al. ("
+2002) and Albert et al. (,2002) and Albert et al. (
+2007). based on,"2007), based on"
+accretion Lows.,accretion flows.
+ As loss-cone or other kinetic instabilities cannot develop in accreting svstems. Ls Cet is not expected o be a maser source.," As loss-cone or other kinetic instabilities cannot develop in accreting systems, ES Cet is not expected to be a maser source."
+ The presence of high-density material would imply a high plasma-cutoll frequency., The presence of high-density material would imply a high plasma-cutoff frequency.
+ For a plasma with electron. number density of ~5lottem 7. the asma frequency will be well above 5 Cllz. thus it will event the propagation of 6 em radio emission. which is he observational band of our ATCA observation.," For a plasma with electron number density of $\sim 5 \times
+10^{11}$ $^{-3}$, the plasma frequency will be well above 5 GHz, thus it will prevent the propagation of 6 cm radio emission, which is the observational band of our ATCA observation."
+ Our observations do not show evidence of radio emission rom RN J1914]24 at a limit of 42 µν., Our observations do not show evidence of radio emission from RX J1914+24 at a limit of 42 $\mu$ Jy.
+ One. obvious »ossibilitv is that Wh may not. occur. in this. svstem., One obvious possibility is that UI may not occur in this system.
+ llowever. the non-detection does not rule UL out.," However, the non-detection does not rule UI out."
+ As »onted. out. in. Willes Wu (2004) the observability of clectron-cyvclotron masers from a UL double-degenerate compact binary depends on the magnetic moment of the magnetic white chwarl (which determines the frequencies of the evelotron harmonics). the amount of thermal electrons lilling the electric-current. Dlowing magnetic Dux tubes. the temperature of these thermal electrons. and the viewing orientation of the binary.," As pointed out in Willes Wu (2004) the observability of electron-cyclotron masers from a UI double-degenerate compact binary depends on the magnetic moment of the magnetic white dwarf (which determines the frequencies of the cyclotron harmonics), the amount of thermal electrons filling the electric-current flowing magnetic flux tubes, the temperature of these thermal electrons, and the viewing orientation of the binary."
+ Caleulations showed that the racio emission is detectable in à restrictive region in the parameter space of Ul double-degenerate compact binaries., Calculations showed that the radio emission is detectable in a restrictive region in the parameter space of UI double-degenerate compact binaries.
+ Thus. even if UL is operating efficiently and clectron-evelotron niasers are generated in all svstems in an ensemble. some syvstenis will show detectable elecetron-evelotron masers in the radio wavebands. while a significant [fraction of the svstems will show null detection in a radio survey.," Thus, even if UI is operating efficiently and electron-cyclotron masers are generated in all systems in an ensemble, some systems will show detectable electron-cyclotron masers in the radio wavebands, while a significant fraction of the systems will show null detection in a radio survey."
+ We have detected a radio source at a position coincident with the known optical position of RA JOSOG|15., We have detected a radio source at a position coincident with the known optical position of RX J0806+15.
+ Although we cannot completely exclude that this is a chance alignment between the known position of RA JOSOG|15 and an artifact in the data reduction process. the fact that it was detected al a significance level of 5.80 and that the radio source was variable suggests that it is more likely that RA 0806|15 is a transient radio source.," Although we cannot completely exclude that this is a chance alignment between the known position of RX J0806+15 and an artifact in the data reduction process, the fact that it was detected at a significance level of $\sigma$ and that the radio source was variable suggests that it is more likely that RX 0806+15 is a transient radio source."
+" With these caveats in münd. we can determine the brightness temperature Z5, of a source by: The distance of RX JOSOG|15 is not well known. with Israel et al (2003)3) noting that the distance to the edge of the Galaxy is 500 pe. while Barros et al (2007) estimate that its distance is greater than 1.1 kpc implying it is out of the Galactic plane."," With these caveats in mind, we can determine the brightness temperature $T_{\rm b}$ of a source by: The distance of RX J0806+15 is not well known, with Israel et al (2003) noting that the distance to the edge of the Galaxy is 500 pc, while Barros et al (2007) estimate that its distance is greater than 1.1 kpc implying it is out of the Galactic plane."
+" Assuming a conservative distance of 500 pe and that the size of the emission region is 2 cm (the linear extension of the foot-point. Dux-tube for systems with a non-magnetic white dwarf companion with mass 0.5 M... see. Willes Wu 2004). the observed. [ux density of 99p.]y implies Zi,=2.1.LOA IK. There are some uncertainties about the exact size of the foot-point emission region."," Assuming a conservative distance of 500 pc and that the size of the emission region is $\sim 2 \times 10^7$ cm (the linear extension of the foot-point flux-tube for systems with a non-magnetic white dwarf companion with mass 0.5 $_\odot$, see Willes Wu 2004), the observed flux density of $\mu$ Jy implies $T_{\rm
+ b}=2.1\times10^{18}$ K. There are some uncertainties about the exact size of the foot-point emission region."
+" Even if we assume that the size of the emission region is IO"" cm. (the radius of a 0.5 M. white dwarf). we still obtain a very high brightness temperature. Z5,—8 LOMA. Such a high brightness temperature cannot be explained by non-thermal svnchrotron process. which would be limited to LOM WS (Dulk Marsh. 1982)."," Even if we assume that the size of the emission region is $10^9$ cm, (the radius of a 0.5 $_\odot$ white dwarf), we still obtain a very high brightness temperature, $T_{\rm b}=8\times10^{14}$ K. Such a high brightness temperature cannot be explained by non-thermal synchrotron process, which would be limited to $<10^{10}$ K (Dulk Marsh 1982)."
+ Ht also cannot be explained by any incoherent radiation processes. as they are limited to ~LOI by inverse Compton cooling (Ixellerman DPaulinv-Loth. 1969).," It also cannot be explained by any incoherent radiation processes, as they are limited to $\sim 10^{12}$ K by inverse Compton cooling (Kellerman Pauliny-Toth 1969)."
+" '""herefore. the radio emission must be generated by a coherent radiative process. such as an electron-cyclotron maser as predicted by the unipolar-induction model (Wu et al."," Therefore, the radio emission must be generated by a coherent radiative process, such as an electron-cyclotron maser as predicted by the unipolar-induction model (Wu et al."
+ 2002)., 2002).
+ In the maser model described in Willes Wu (2004). the transient or bursting nature of the source may be explained by variations in the emission-cone beaming direction or by the presence of a small amount. of non-thermal electrons whose densitv Ductuate.," In the maser model described in Willes Wu (2004), the transient or bursting nature of the source may be explained by variations in the emission-cone beaming direction or by the presence of a small amount of non-thermal electrons whose density fluctuate."
+ To confirm the nature of the radio source we urge further observations of this source at radio wavelengths to determine how often it shows radio emission and to better constrain the upper limit on the circular polarisation., To confirm the nature of the radio source we urge further observations of this source at radio wavelengths to determine how often it shows radio emission and to better constrain the upper limit on the circular polarisation.
+ The fact that we did not detect any evidence for racio emission from C.J S76 does not rule out the operation of UL in this svstem., The fact that we did not detect any evidence for radio emission from GJ 876 does not rule out the operation of UI in this system.
+ Our observations of C.J ST6 took place at 12 mn., Our observations of GJ 876 took place at 12 mm.
+ The caleulations of Willes Wu (2005) suggest that any racio emission due to the Ul operation would more likely be observable at shorter wavelengths., The calculations of Willes Wu (2005) suggest that any radio emission due to the UI operation would more likely be observable at shorter wavelengths.
+ Sensitive observations at these wavelengths (<6 mm) will be possible using ALMA., Sensitive observations at these wavelengths $<$ 6 mm) will be possible using ALMA.
+ The Australia Telescope is funded by the Commonwealth of Australia for operation as à National Facility managed. by the CSIRO., The Australia Telescope is funded by the Commonwealth of Australia for operation as a National Facility managed by the CSIRO.
+ The National Raclio Astronomy Observatory is a facility of the National Science. Foundation of the USA. operated under cooperative agreement by Associated Universities. Inc. We thank the stall of both the AT ancl VLA for assistance with the observations and Steven Longmore for help in reducing the ES Cet cata.," The National Radio Astronomy Observatory is a facility of the National Science Foundation of the USA, operated under cooperative agreement by Associated Universities, Inc. We thank the staff of both the AT and VLA for assistance with the observations and Steven Longmore for help in reducing the ES Cet data."
+ INN thanks Richard Lunsteacl for. discussions., KW thanks Richard Hunstead for discussions.
+ Armagh Observatory is erant aided by the N. Ireland. Dept., Armagh Observatory is grant aided by the N. Ireland Dept.
+ of Culture. Arts and Leisure.," of Culture, Arts and Leisure."
+ , 	 
+lime / [rom 2—x by ,time $t$ from $z\rightarrow\infty$ by gt.
+We apply this initial condition to the particle model., We apply this initial condition to the particle model.
+ An equivalent condition is that the comoving coordinate displacements [rom (he particle positions at /=0 scale as αν , An equivalent condition is that the comoving coordinate displacements from the particle positions at $t=0$ scale as $\vec x(t)-\vec x(0)\propto a(t)$ .
+This is an approximation. among other reasons because it does not Cake account of the nonlinear development of structure on the scale of galaxies αἱ ο.," This is an approximation, among other reasons because it does not take account of the nonlinear development of structure on the scale of galaxies at $z_i$."
+ The point is that it fits the idea that peculiar velocities of protogalaxies al high recdshilt are small ancl growing as structure grows., The point is that it fits the idea that peculiar velocities of protogalaxies at high redshift are small and growing as structure grows.
+ The third approximation is the use of four external actors lor a phenomenological description of the effect of matter external to LG on the motions of the LG galaxies., The third approximation is the use of four external actors for a phenomenological description of the effect of matter external to LG on the motions of the LG galaxies.
+ The actors represent in first approximation the neighboring mass concentrations in the Sculptor Group and the Madfei and Centaurus svstems. but the masses. positions and redshifts of the actors are allowed to float to improve the model fit to the measurements of the LG ealaxies.," The actors represent in first approximation the neighboring mass concentrations in the Sculptor Group and the Maffei and Centaurus systems, but the masses, positions and redshifts of the actors are allowed to float to improve the model fit to the measurements of the LG galaxies."
+ The approximation would fail if a large external mass that now has little effect on LG was a serious influence in the past. when these masses were much closer to LG.," The approximation would fail if a large external mass that now has little effect on LG was a serious influence in the past, when these masses were much closer to LG."
+ This need not be a problem. of course. because LG was more compact too.," This need not be a problem, of course, because LG was more compact too."
+ And we have a meaningful indication that this and the other (wo approximations vield a useful approach to reality. from the generally reasonable fit of the model to considerably more measurements than there are relevant. adjustable parameters.," And we have a meaningful indication that this and the other two approximations yield a useful approach to reality, from the generally reasonable fit of the model to considerably more measurements than there are relevant adjustable parameters."
+ But an important goal for a more complete analvsis is (o check the third approximation by taking explicit account of the dvnanmies and gravitational effects of the observed external mass concentrations., But an important goal for a more complete analysis is to check the third approximation by taking explicit account of the dynamics and gravitational effects of the observed external mass concentrations.
+ In the form applied here the numerical action method (NAM) of dealing with the mixed boundary conditions of given. present positions and the initial condition in equation (2.1)) represents particle orbits by positions at discrete time steps (Peebles 1989. 1995: Peebles. Phelps. 5hava. Tully 2001).," In the form applied here the numerical action method (NAM) of dealing with the mixed boundary conditions of given present positions and the initial condition in equation \ref{eq:gt}) ) represents particle orbits by positions at discrete time steps (Peebles 1989, 1995; Peebles, Phelps, Shaya, Tully 2001)."
+ It produces a solution to the equation of motion in leapfrosg approximation that is reached by iterated joint shifts of all particle positions from random (rial orbits toward a stationary point of the action in the direction indicated by (he first aud second derivatives of the action., It produces a solution to the equation of motion in leapfrog approximation that is reached by iterated joint shifts of all particle positions from random trial orbits toward a stationary point of the action in the direction indicated by the first and second derivatives of the action.
+ The method is described in Appendix A., The method is described in Appendix A.
+Iverocdvnamical simulations of the cireiumstellar mecditm of low-to-intermediate mass stars (LIMS) have traditionally focused on the morphology of the gas and neglected. the presence of dust erains in the stellar wind (E.g.Wareingοἱal.2007).,Hydrodynamical simulations of the circumstellar medium of low-to-intermediate mass stars (LIMS) have traditionally focused on the morphology of the gas and neglected the presence of dust grains in the stellar wind \citep[E.g.][]{wetal07}.
+. Those simulations that do include dust tend to focus on the inner wind region (Woitke2006) or deal with the proto-planetary disks surrounding verv voung stars (Paardekooper&Mellema2006)., Those simulations that do include dust tend to focus on the inner wind region \citep{w06} or deal with the proto-planetary disks surrounding very young stars \citep{pm06}.
+. With recent aclvances in infrared observations. such as satellites like Spitzer ancl LIerschel. it has become possible to make detailed infrared observations of the circumstellar environment.," With recent advances in infrared observations, such as satellites like Spitzer and Herschel, it has become possible to make detailed infrared observations of the circumstellar environment."
+ The higher sensitivity of these satellites has enabled us to observe the weak signatures of the bowshocks that form when the wind ofa moving star collides with the surrounding interstellar medium (ISM)., The higher sensitivity of these satellites has enabled us to observe the weak signatures of the bowshocks that form when the wind of a moving star collides with the surrounding interstellar medium (ISM).
+ Since dust erains are the primary source of infrared. radiation. it becomes necessary to fully integrate them into the simulations to determine how the distribution of dust grains correlates with the morphology of the circumstellar eas aud how the presence of dust influences the structure of the shock that occurs where the stellar wind interacts with ihe ISM.," Since dust grains are the primary source of infrared radiation, it becomes necessary to fully integrate them into the simulations to determine how the distribution of dust grains correlates with the morphology of the circumstellar gas and how the presence of dust influences the structure of the shock that occurs where the stellar wind interacts with the ISM."
+ As an example. we have investigated the behavior of dust erains in the cireumstellar environment of a fast moving red supereiant star.," As an example, we have investigated the behavior of dust grains in the circumstellar environment of a fast moving red supergiant star."
+ As input parameters we use the values obtained fora-Orionis.. which has been observed in infrared with the AINXARI Satellite 2008).," As input parameters we use the values obtained for, which has been observed in infrared with the AKARI Satellite \citep{uetal08}."
+. This star shows evidence of a bow shock. indicating that it is moving through the ISM at à velocity of 40ng kms with n; the dimensionless local hvdrogen clensity 2008).," This star shows evidence of a bow shock, indicating that it is moving through the ISM at a velocity of $40\,n_H^{-0.5}$ km/s with $n_H$ the dimensionless local hydrogen density \citep{uetal08}."
+. 50 lar. most simulations of fast-moving stars have loceused on hot. high mass stars (BrighenG&DErcole1995:ComeronKaper1998:vanAMarleetal. 2006).. while simulations of the bowshocks of cool. lower mass stars were done bv. Wareingetal.(2007).," So far, most simulations of fast-moving stars have focused on hot, high mass stars \citep{be95,ck98,vmetal06}, while simulations of the bowshocks of cool, lower mass stars were done by \citet{wetal07}."
+. llowever. none of these simulations took the presence of dust into account.," However, none of these simulations took the presence of dust into account."
+ llf one wants to incorporate dustgrains in the hydrodvnamical models. one should not only take into account the motion of the dustgrains themselves. it also becomes necessary {ο add the kinetic interaction between eas and dust to the equations of hydrodynamics.," If one wants to incorporate dustgrains in the hydrodynamical models, one should not only take into account the motion of the dustgrains themselves, it also becomes necessary to add the kinetic interaction between gas and dust to the equations of hydrodynamics."
+MOND inodel is a τσ deviation.,MOND model is a $7\sigma$ deviation.
+ We found a better solution (rj=30 kpe aud other parineters the same}. which places MOND *unl at 6.20.," We found a better solution $r_b=30$ kpc and other parameters the same), which places MOND “only” at $6.2\sigma$."
+" Tn order to cuvestigate wiat stellar ass ds receded. for MOND. we split the sample of red omnumrnes mto 5 mass bius remeine from AM.~5&10!AF, to AL~5«10!AL..."," In order to envestigate what stellar mass is needed for MOND, we split the sample of red primaries into 5 mass bins ranging from $M_*\sim
+5\times 10^{10}M_\odot$ to $M_*\sim 5\times 10^{11}M_\odot$."
+ We chose the ess stringent isolation condition because it gives vpically 2-2.5Γ time more sacllites resulting in sunaller statistical errors., We chose the less stringent isolation condition because it gives typically 2-2.5 time more satellites resulting in smaller statistical errors.
+ Eacth bin has a large nunhber of satellites: 12002900., Each bin has a large number of satellites: $1200-2900$.
+ We then runi a erid of MOND inodels wihn different asses and select models. which make best fits.," We then run a grid of MOND models with different masses and select models, which make best fits."
+ Just as the MOND inodels in Figure L.. there are models. which mareinally fit the data.," Just as the MOND models in Figure \ref{fig:Tuned}, there are models, which marginally fit the data."
+ For each model we ect stellar mass required bv MOND and compare it with the stellar mass estimated bv the stellar population models, For each model we get stellar mass required by MOND and compare it with the stellar mass estimated by the stellar population models.
+ Results are preseuted in Figure («, Results are presented in Figure \ref{fig:MONDmass}.
+ For smaller 11asses MOND eives masses. Which are compatible with actual stellar mass observed in the galaxies.," For smaller masses MOND gives masses, which are compatible with actual stellar mass observed in the galaxies."
+ The plot clearly shows the problem with massive galaxies: MOND requires increasingly lore πας than observed in the galaxies cuding up in large ( 2.5) disagreement with the observatious., The plot clearly shows the problem with massive galaxies: MOND requires increasingly more mass than observed in the galaxies ending up in large $\sim 2.5$ ) disagreement with the observations.
+ The external eravity force ger ds another MOND compoucut., The external gravity force $g_{\rm ext}$ is another MOND component.
+ It iust exist on the level of Yor09OOlay.," It must exist on the level of $g_{\rm ext}\approx
+0.01a_0$."
+ Figure 7 shows what lappeus to the models when we add the external force., Figure \ref{fig:TunedExt} shows what happens to the models when we add the external force.
+ As tle starting models we use the best fits shown in the Figure Las dashed curves., As the starting models we use the best fits shown in the Figure \ref{fig:Tuned} as dashed curves.
+ The models with the external force make the fits wach worse., The models with the external force make the fits much worse.
+ Taken at face value. the models with realistic gar=0.015609 can be rejected.," Taken at face value, the models with realistic $g_{\rm ext}=0.015a_0$ can be rejected."
+ Again. we can make the model work if we inercase ML hy a factor ~1.5.," Again, we can make the model work if we increase M/L by a factor $\sim
+1.5$."
+ Yet. this will make the situation with stellar masses even worse. than what we already have.," Yet, this will make the situation with stellar masses even worse, than what we already have."
+ Using the SDSS DRI data we study the distribution aud velocities of satellites orbiting red isolated galaxies., Using the SDSS DR4 data we study the distribution and velocities of satellites orbiting red isolated galaxies.
+ Wefind that the surface nuuber-deusitv of he satellites declines almost as à power law with the slope 2.53., Wefind that the surface number-density of the satellites declines almost as a power law with the slope $-2.5-3$.
+ The distribution of the line-ofsight velocities is nearly a perfect Gaussian distribution with a coustant component due to interlopers., The distribution of the line-of-sight velocities is nearly a perfect Gaussian distribution with a constant component due to interlopers.
+ The nus velocities are found to eradually decline with the projected cistauce., The rms velocities are found to gradually decline with the projected distance.
+ The coustaut rius velocity (isothermal solution) can be rejected at a 10 σ level., The constant rms velocity (isothermal solution) can be rejected at a 10 $\sigma$ level.
+ Observational data strongly favor the standard cosinological model: all three major statistics of satellites — the. nuniber-deusitv profile. the liue-ofsight velocity dispersion. a1 the distribution function of the velocities agre! remarkably well with the theoretical predictions.," Observational data strongly favor the standard cosmological model: all three major statistics of satellites – the number-density profile, the line-of-sight velocity dispersion, and the distribution function of the velocities – agree remarkably well with the theoretical predictions."
+ Thus. the success of the standard model extens fo scales (40-500 kpe. much lower than what was previously considered.," Thus, the success of the standard model extends to scales (50-500) kpc, much lower than what was previously considered."
+ MOND fails badly in cases with any acceptable »ower- law approximation for the mmuber-density of satellites aud a constant velocity anisotropy by xoducimg sharply declining velocities at small distances followed by nearly fiat velocities at aree distances — just the opposite of what is observed iu real galaxies., MOND fails badly in cases with any acceptable power- law approximation for the number-density of satellites and a constant velocity anisotropy by producing sharply declining velocities at small distances followed by nearly flat velocities at large distances – just the opposite of what is observed in real galaxies.
+ Models way be made to Bt the satellites data ouly when all the following conditious are fulfilled: We fud that the later three conditions are difficult to realize in nature., Models may be made to fit the satellites data only when all the following conditions are fulfilled: We find that the later three conditions are difficult to realize in nature.
+ Satellites do not fall in to their parent galaxies with zero tanecutial velocities., Satellites do not fall in to their parent galaxies with zero tangential velocities.
+ The velocities are induced by other neighboring galaxies aud by large-scale structures such as fluueuts., The velocities are induced by other neighboring galaxies and by large-scale structures such as filaments.
+ To some degree. it is simular to the tidal torque. which is responsible for the oriein of the augular momentum of galaxies.," To some degree, it is similar to the tidal torque, which is responsible for the origin of the angular momentum of galaxies."
+ Yet. the interactions are more efficient for providing random volocities.," Yet, the interactions are more efficient for providing random velocities."
+ Measurements of peculiar velocities of ealaxies the Local Volume (3-5 Mpc around the Milky Wav) found deviations frou the IIubble flow that about 70-80 kin/s (Macció 2005).., Measurements of peculiar velocities of galaxies the Local Volume (3-5 Mpc around the Milky Way) found deviations from the Hubble flow that about 70-80 km/s \citep{Maccio2005}. .
+ At these distances the eravitational pull of other large galaxies iu the aria is larger, At these distances the gravitational pull of other large galaxies in the aria is larger
+alternative idea was proposed by Pringle(1989).. who sugeested that clusters can be heated by dissipation of sound waves generated by galaxy motions in the cluster.,"alternative idea was proposed by \citet{pr89}, who suggested that clusters can be heated by dissipation of sound waves generated by galaxy motions in the cluster."
+ Further support for the idea that viscosity may play au Huportant role iu the intracluster imnediunm comes frou the recent study of density profiles in clusters (ITanseu&Stadel 2003)., Further support for the idea that viscosity may play an important role in the intracluster medium comes from the recent study of density profiles in clusters \citep{hs03}.
+.. The main purpose of this paper is to demonstrate that clusters can be heated oeffücieutlv by wave dissipation associated with activity of ACN located in their centers., The main purpose of this paper is to demonstrate that clusters can be heated efficiently by wave dissipation associated with activity of AGN located in their centers.
+ Although our simulations are two-dimensional aud therefore not directly applicable to real clusters. we argue that the basic result should be preserved in three dinieusions.," Although our simulations are two-dimensional and therefore not directly applicable to real clusters, we argue that the basic result should be preserved in three dimensions."
+" The intracluster mecdimm is mmitiallv emet to be iu hydrostatic equilibrimim in an NEW potential acceleration(Navarrectal.1995.1997) for which the eravitational as a function of the distance from the cluster center rk ds eiven bv r,=[00 kpe is the core radius. e=rf/r,.. B=yr. 6.where=Tu3.0«10! is the central overdensity. aud poAMT.D is the critical deusity of the Universe (we assume ZZ,το lan ! +)."," The intracluster medium is initially assumed to be in hydrostatic equilibrium in an NFW potential \citep{na95,na97} for which the gravitational acceleration as a function of the distance from the cluster center $r$ is given by where $r_{c}=100$ kpc is the core radius, $x=r/r_{c}$, $\mathbf{\hat{r}}=\mathbf{r}/r$, $\delta_c = 3.0\times 10^{4}$ is the central overdensity, and $\rho_{\rm crit}=3H_{o}^{2}/(8\pi G)$ is the critical density of the Universe (we assume $H_{o}=75$ km $^{-1}$ $^{-1}$ )."
+ The initial teniperature distribution is given by where T=3.0 keV. 10 kpc aud ο=0.22.," The initial temperature distribution is given by where $T_{o}=3.0$ keV, $r_{o}=10$ kpc and $\beta =0.22$."
+ The teiiperature at 100 kpe is 5.1 keV. The central electron muuber density is 2.8«10.2 c.7.," The temperature at 100 kpc is $5.1$ keV. The central electron number density is $2.8\times
+10^{-2}$ $^{-3}$."
+ The electron ποτ deusity at LOO kpe is approximately 5.3.10P 7., The electron number density at 100 kpc is approximately $5.3\times 10^{-3}$ $^{-3}$.
+ This corresponds to a ceutral cooling time of ~1.3«10° vers and a cooling time of ~1.0&1019 years at 100 kpe.," This corresponds to a central cooling time of $\sim 1.3\times 10^{9}$ years and a cooling time of $\sim 1.0\times
+10^{10}$ years at 100 kpc."
+ These are values typical of a cooling How cluster as the central cooling time is wich shorter than the Iubble time., These are values typical of a cooling flow cluster as the central cooling time is much shorter than the Hubble time.
+ Note that we do not imtend to model the details of any particular cluster even if our studies have Όσοι motivated by observations of Perseus., Note that we do not intend to model the details of any particular cluster even if our studies have been motivated by observations of Perseus.
+ Since our simulations are two-dimensional. a quantitative comparison is not possible.," Since our simulations are two-dimensional, a quantitative comparison is not possible."
+ For oue thing. the wave eneregv deusity can be expected to decay more slowly with radius in two dimensions (xr. 1) than in three (x 7).," For one thing, the wave energy density can be expected to decay more slowly with radius in two dimensions $\propto
+r^{-1}$ ) than in three $\propto r^{-2}$ )."
+" Moreover. bubbles expanding iuto dimensions are expected tobe larger than equivaleut wwbubbles expanding three-dinenusionallv: iudeced. for our initial conditions chosen to those in Perseus. we find that our computed bubbles ""prestanteare larger than the observed cavities."," Moreover, bubbles expanding into two dimensions are expected to be larger than equivalent bubbles expanding three-dimensionally; indeed, for our initial conditions chosen to approximate those in Perseus, we find that our computed bubbles are larger than the observed cavities."
+ We asstune that the eas is fully ionized aud characterized by X=0.75 aud Y=0.25. where X aud Y are the lvdrogen aud helm fractions.," We assume that the gas is fully ionized and characterized by $X=0.75$ and $Y=0.25$, where $X$ and $Y$ are the hydrogen and helium fractions."
+" The injected gas is characterized by an acliabatic iudex panne= 1/3. whereas for the ambient eas we used ryan,= levels5/3."," The injected gas is characterized by an adiabatic index $\gamma_{\rm bubble}=4/3$ , whereas for the ambient gas we used $\gamma_{\rm
+ICM}=5/3$ ."
+ Calculations were done in two dimensions for 9 of refinement usine the PPAL adaptive mesh refineumieut code FLASIL, Calculations were done in two dimensions for 9 levels of refinement using the PPM adaptive mesh refinement code FLASH.
+ The size of the computational domain was (200 ο)”., The size of the computational domain was (200 $^2$.
+ Thus. the effective resolution im our siumlatious was 20187? zones. which correspouds to ~0.1 spe.," Thus, the effective resolution in our simulations was $^2$ zones, which corresponds to $\sim 0.1$ kpc."
+ We have performed vouvergenee tests and found hat neither the geometry of the bubbles nor that of he waves depeuds on the adopted resolution., We have performed convergence tests and found that neither the geometry of the bubbles nor that of the waves depends on the adopted resolution.
+ This is due to the fact that. for the parameters considered. the Revuolds umuber corresponding to bubbles aud waves is ow and also because numerical dissipation of PPM codes is known to be relatively low.," This is due to the fact that, for the parameters considered, the Reynolds number corresponding to bubbles and waves is low and also because numerical dissipation of PPM codes is known to be relatively low."
+ The effective. Revuolds imnbers achievable iu the simulation are proportional o the nmuuber of grid poiuts across the fluctuation of interest to the power η. where s=3 is the order of the umuerical (Sytineetal.2000:Bal-musctal.1996:Porter&Woodward 1991).," The effective Reynolds numbers achievable in the simulation are proportional to the number of grid points across the fluctuation of interest to the power $n$, where $n=3$ is the order of the numerical \citep{syt00,bal96,por94}."
+. We used a redefined system of umts iu which alb variables apart from temperature are close to unity and adopted outflow boundary couditious., We used a redefined system of units in which all variables apart from temperature are close to unity and adopted outflow boundary conditions.
+ We lave also repeated our Sa.ulations in two-dimensional cevliudrical coordinates and found that. in spite of uuphlivsical effects near the svuuuetryv axis. our conclusions remain unaffected.," We have also repeated our simulations in two-dimensional cylindrical coordinates and found that, in spite of unphysical effects near the symmetry axis, our conclusions remain unaffected."
+ We model AGNheating bv injecting hot eas iuto two regions of radius 1 kpc located 10 kpc to either side of the cluster ceuter., We model AGN heating by injecting hot gas into two regions of radius 1 kpc located 10 kpc to either side of the cluster center.
+ The energyinjection rate. £. for cach source and the mass injection rate per unit volume. p. are both constant.," The energyinjection rate, $L$, for each source and the mass injection rate per unit volume, $\dot{\rho}$, are both constant."
+ Thus. the energy injection rate per unit mass é is computed from where V ds the volume of one injection region (of radius l kpe).," Thus, the energy injection rate per unit mass $\dot{\epsilon}$ is computed from where $V$ is the volume of one injection region (of radius 1 kpc)."
+ We used £=L5«10! cre ! aud pV.=0.01 AL.sr+., We used $L=1.5\times 10^{43}$ erg $^{-1}$ and $\dot{\rho}V=0.01$ $_{\odot}$ $^{-1}$.
+ The energy injection is intermittent with au intermittency period of 3«10* years. Lie. the source is active for 1.5«10* yearsaud dormant for 1.5«107 vears.," The energy injection is intermittent with an intermittency period of $3\times 10^{7}$ years, i.e., the source is active for $1.5\times 10^{7}$ yearsand dormant for $1.5\times 10^{7}$ years."
+" In the initial state for each activity episode. the teiiperature and deusitv are a hundred times higher aud lower. respectively, than the temperature aud deusitv in the initial unperturbed state at the same The dissipation of mechanical energy. due to viscosity. per unit mass of the fuid. was calculated from where A=¢;; aud and where gis the dynamical coefficientofviscosity."," In the initial state for each activity episode, the temperature and density are a hundred times higher and lower, respectively, than the temperature and density in the initial unperturbed state at the same The dissipation of mechanical energy due to viscosity, per unit mass of the fluid, was calculated from \citep{bat67,shu92}
+ where $\Delta =e_{ii}$ and and where $\mu$ is the dynamical coefficientofviscosity."
+ We use the standard Spitzer viscosity (Bragiuskii LO58).. for which p=1.1<10MT? e ts ," We use the standard Spitzer viscosity \citep{bra58}, , for which $\mu = 1.1\times
+10^{-16}T^{5/2}$ g $^{-1}$ $^{-1}$ ."
+As conditious inside thebuovautlv rising Ijbles are very uncertain and because we want to focus on enerev dissipation in, As conditions inside the buoyantly rising bubbles are very uncertain and because we want to focus on energy dissipation in
+by h(r.f),"by $h(\vec r,t)$."
+ In short. pulse propagation through complete eveles of hf) have uo net effect ou the arrival time.," In short, pulse propagation through complete cycles of $h(\vec r,t)$ have no net effect on the arrival time."
+ There is oulv a perturbation of the arrival time by. the incomplete cycles traversed at the pulsar aud at Earth., There is only a perturbation of the arrival time by the incomplete cycles traversed at the pulsar and at Earth.
+ We use Eqn., We use Eqn.
+ 16 in PAI63 to caleulate the luuinosity Leay of the gravitational wave (CAV) from the assumed circular binary svstem in the GC where s)=2.6&105 is the total mass of the Chezetal. 2000).. 4 is the mass ratio (q<1). aud a is the seuniauajor axis.," 16 in PM63 to calculate the luminosity $L_{GW}$ of the gravitational wave (GW) from the assumed circular binary system in the GC where $m=2.6\times 10^6$ is the total mass of the \cite[]{Ghez00}, , $q$ is the mass ratio $q\le 1$ ), and $a$ is the semi-major axis."
+ The uuuerical result uses Iepler's Law «=Gnta? and the 106d. orbital period of interest., The numerical result uses Kepler's Law $a^3=GmP_{orb}^2/4\pi^2$ and the 106d orbital period of interest.
+" For Po,Pl=1064.4 =59AU."," For $P_{orb}=106d$, $a=59{\rm~AU}$."
+ Tere aud. below we assunue a circular orbit which is Likely following the combined actions of ανασα. friction aud radiation., Here and below we assume a circular orbit which is likely following the combined actions of dynamical friction and radiation.
+ Frou. energv deusitv of a CAV. which is (=sanplitude25243926' the‘of(Equ.," From the energy density of a GW, which is $U = c^2\dot{h}^2/32\pi G$ (Eqn."
+ 2 in PM623). we derive the dimensionless the CW: where d ds the distance to the cimitter. 8 kpe.," 2 in PM63), we derive the dimensionless amplitude of the GW; where $d$ is the distance to the emitter, 8 kpc."
+ In this expression 7) ds averaged over alb orientations of the observer rolative to the plane of the binary orbit., In this expression $h$ is averaged over all orientations of the observer relative to the plane of the binary orbit.
+ Since we obtained this expression from the total energy density. which is the sum of the coutributions from two polarizations. ερ aud Ji.. fis actually the quadrature sua of these two polarizations.," Since we obtained this expression from the total energy density, which is the sum of the contributions from two polarizations, $h_+$ and $h_\times$, $h$ is actually the quadrature sum of these two polarizations."
+ Iu order to find the dependence ou iuclination angle. /. we use the expression for the average power radiated per solid angle in PAIG3.," In order to find the dependence on inclination angle, $i$, we use the expression for the average power radiated per solid angle in PM63."
+ This shows that the power radiated along the axis of the orbit Is 5 times that for an edee-on view., This shows that the power radiated along the axis of the orbit is 8 times that for an edge-on view.
+" As discussed. by. Detweiler(1979) aud. others. the dineusiouless strain / produces au apparent redshift iu the pulsu frequency,", As discussed by \cite{Detweiler79} and others the dimensionless strain $h$ produces an apparent redshift in the pulsar frequency.
+" À periodic source of CW then will produce a periodic shift iu pulse arrival time from propagation through the eravitational radiation with an amplitude. f. which is given by where Poy represents the period of the eravitational wave. P,,7/2."," A periodic source of GW then will produce a periodic shift in pulse arrival time from propagation through the gravitational radiation with an amplitude, $\delta t$, which is given by where $P_{\rm GW}$ represents the period of the gravitational wave, $P_{orb}/2$."
+ The mass ratio factor is at most 1/1., The mass ratio factor is at most $1/4$.
+ The angular factor ranges from 1 to 2.8., The angular factor ranges from 1 to 2.8.
+ Therefore. df is less than 16 us. aud its average value over all solid angles is which is us for g=1 aud less for iuv other mass ratio.," Therefore, $\delta t$ is less than 16 ns, and its average value over all solid angles is which is 11 ns for $q=1$ and less for any other mass ratio."
+ Detweiler(1979). discusses the dependence of the GV sjenature in pulsar timing on the augle between the CAV and the pulsar sighthue., \cite{Detweiler79} discusses the dependence of the GW signature in pulsar timing on the angle between the GW and the pulsar sightline.
+ Pulse propagation times are perturbed bv the CAV owing to incomplete traversal of a cvcle of the GW both at the pulsar as pulses are emitted and at the Earth upon pulse reception., Pulse propagation times are perturbed by the GW owing to incomplete traversal of a cycle of the GW both at the pulsar as pulses are emitted and at the Earth upon pulse reception.
+" In the plane wave approxination the resulting timine residual. δὲ, is where 5 is the cosine of the angle. o. between the GC and the pulsar. where o=0 is defined as the pulsar Wing along the line of sight to the GC."," In the plane wave approximation the resulting timing residual, $\delta t$, is where $\gamma$ is the cosine of the angle, $\phi$, between the GC and the pulsar, where $\phi=0$ is defined as the pulsar lying along the line of sight to the GC."
+ F(f) is the dimensionless phase term that comes from the fragments of the CAV traversed at the emitter (ο) aud receiver () ends (f£<1)., $f(t)$ is the dimensionless phase term that comes from the fragments of the GW traversed at the emitter (e) and receiver (r) ends $f\le1$ ).
+" The times of enission and reception are £. aud f;=τς{1ο. respectively, and the factor 5/ is the projection of the pulsar distance / along the CAV propagation vector."," The times of emission and reception are $t_{\rm e}$ and $t_{\rm r}=t_{\rm e}+l/c$, respectively, and the factor $\gamma l$ is the projection of the pulsar distance $l$ along the GW propagation vector."
+ Note that when +=|1. there is no effect from the CAV: a pulsar Iwineg along the line of sight to the GC will experieuce no effect.," Note that when $\gamma=+1$, there is no effect from the GW: a pulsar lying along the line of sight to the GC will experience no effect."
+ The residual is also identically zero for στ=1 which describes electromagnetic waves (EMA) traveling in the opposite direction as the CAV., The residual is also identically zero for $\gamma=-1$ which describes electromagnetic waves (EMW) traveling in the opposite direction as the GW.
+ The residual iucreases as the EADW and CAV become perpendicular. aud reaches a maxinumn just before they become parallel.," The residual increases as the EMW and GW become perpendicular, and reaches a maximum just before they become parallel."
+ The augle between PSR B1937|21 and Ser A from Earth is and between PSR J17123|0717 aud Ser is.," The angle between PSR B1937+21 and Sgr $^*$ from Earth is, and between PSR J1713+0747 and Sgr $^*$ is."
+. The sienal from the two pulsars will therefore be diminished from our carher estimate in Equation Lo bv a factor (1135)/2=0.76 and 0.90. respectively.," The signal from the two pulsars will therefore be diminished from our earlier estimate in Equation \ref{eqn:46ns} by a factor $(1+\gamma)/2 = 0.76$ and 0.90, respectively."
+ This factor adjusts our earlier estimate of the maximum possible effect. from 16 us to llus aud our estimate of the average effect from 11 us to 10 us due to an equalimass binary.," This factor adjusts our earlier estimate of the maximum possible effect from 16 ns to 14 ns, and our estimate of the average effect from 11 ns to 10 ns due to an equal-mass binary."
+ Tfeither of the sources were nuch closer to the GC than the Earth. then  in Equ.," If either of the sources were much closer to the GC than the Earth, then $h$ in Eqn."
+ 5 would need to be corrected for both the cdiffereut wave amplitudes aud the different aneular factors., \ref{eqn:detweiler} would need to be corrected for both the different wave amplitudes and the different angular factors.
+ In our case all relevant distances (Earth to CC. J1712|OT17 to GC. DI937than|21 to GC) are 7-8 kpc. and are not known to better (Ikaspi.Tavlor.&Ryba1991:Camilo.Thorsett.Iulkuni 1991).," In our case all relevant distances (Earth to GC, J1713+0747 to GC, B1937+21 to GC) are 7-8 kpc, and are not known to better than \citep{Kaspi94, Camilo94}."
+. Since the distances aud therefore the phase factors are unknown. and the amplitudes of the two effects at the two cuds are roughly equal. it is even possible that the two phase factors will be nearly the same. vastly diminishing the signal.," Since the distances and therefore the phase factors are unknown, and the amplitudes of the two effects at the two ends are roughly equal, it is even possible that the two phase factors will be nearly the same, vastly diminishing the signal."
+ The signal at cuitter site will represent CW level 104 years ago. probably uot too different from today.," The signal at emitter site will represent GW level $10^4$ years ago, probably not too different from today."
+ The phase factor from the receiver cud will produce a signature that is correlated with other pulsus. whereas the emission ternis will be uncorrelated.," The phase factor from the receiver end will produce a signature that is correlated with other pulsars, whereas the emission terms will be uncorrelated."
+ Clearly observations with an array of pulsars at different angles aud distances are critical to overcome this “emission phase” noise.," Clearly observations with an array of pulsars at different angles and distances are critical to overcoming this “emission phase"" noise."
+ We have used mass ratio of 4=1 to calculate the πανκαπ signal we ο expect., We have used a mass ratio of $q=1$ to calculate the maximum signal we might expect.
+ There are two areuinents against a large mass ratio., There are two arguments against a large mass ratio.
+ One is based ou evolutionary arguments and the other ou proper motion observatious of Ser A., One is based on evolutionary arguments and the other on proper motion observations of Sgr $^*$.
+ First. the lifetime for gravitational inspiral from a 106d orbit for g=1 is 3«109 x: the lifetime increases with (114)2/4.qx1.," First, the lifetime for gravitational inspiral from a 106d orbit for $q=1$ is $3\times 10^6$ y; the lifetime increases with $(1+q)^2/q,~q\le 1$."
+" After two galaxies merece the timescale for their ceutral black holes to reach such an orbit is unknown. aud could be from 30 nuillion vears to a IIubble time (0... Rajagopal Romani 1995. Could Ris 2000 3). Tothi""O"," After two galaxies merge the timescale for their central black holes to reach such an orbit is unknown, and could be from 30 million years to a Hubble time (e.g., Rajagopal Romani 1995, Gould Rix 2000 \nocite{Gould00, Rajagopal95}) )."
+striker(1992) rule out the possibility of a recent merger using models of disk heating via accretion of satellite galaxies., \cite{Toth92} rule out the possibility of a recent merger using models of disk heating via accretion of satellite galaxies.
+" They demonstrate that uo more than the mass of the ealaxy could have been accreted within the last 5 billion ταν,", They demonstrate that no more than of the mass of the galaxy could have been accreted within the last 5 billion years.
+ Na&Os-model the acemmulation of a central \IBIT in a scenario with much less disk beating: the accretion of primordial —109ML. black holes., \cite{Xu94} model the accumulation of a central MBH in a scenario with much less disk heating: the accretion of primordial $\sim 10^6~\Msolarm$ black holes.
+ They show that a quickly acciuulatiug ΕΟΤ in the GC is actually not what we should expect: dynamical friction and gravitational radiation serve to eject massive objects from the center, They show that a quickly accumulating MBH in the GC is actually not what we should expect; dynamical friction and gravitational radiation serve to eject massive objects from the center
+"where Mi, is the total. virial mass of the cluster. fyas is the hot gas mass fraction of clusters and D,(2) is the augularsize distance.","where $M_{tot}$ is the total, virial mass of the cluster, $f_{gas}$ is the hot gas mass fraction of clusters and $D_a(z)$ is the angular–size distance."
+ Iu the last line. we have used the fact that there exists a - . T ↑↕∶↴∙⊾∐↑↥⋅↸∖↕⋜↧↑↕∪∐↴⋝↸∖↑↖↖⇁↸∖↸∖∐⊸∖↥⋅⋜↧⋅↖↽↑↸∖∐∏⋉∖↥⋅⋜↧⊓∐⋅↸∖⋜⋯≼↧↖↽∐⋅↕⋜↧↕⋯⋜↧↴∖∷∖↴∶⊺∖⋀∐−⊽⋟≺↓↖−∙⋝," In the last line, we have used the fact that there exists a tight relation between X–ray temperatureand virial mass: $T \propto M^{2/3} (1+z)$ (Evrard et al."
+≺⊏∏⋅⋜∐⋅, 1996).
+"≼⊓∖↑⋜↧↕∙↓∩≝⋊≱⋝⋅∫⇀↸∖↑↴∖↴ fo Dm+ +⋅ ⋅ ↸⊳∪∐∏≻⋜∐⋅↸∖↑↕∐↴∖↴↖↖↽↕↑∐↑∐↸∖↸⊳∪↥⋅↥⋅↸∖↴∖↴⋯∐≼∐∐∶↴⋁↸∖⊼↻↥⋅↸∖↴∖∷∖↴↕∪∐↕⋟∪↥⋅↑∐↸∖⊸∖⊽↥⋅⋜↧⋅↖↽∏∏⊼∪↕⋟⋜↧↸⊳↕∏↴∖↴↑↸∖↥⋅∶ with D, denoting the huninosity distance.", Let's compare this with the corresponding expression for the X–ray flux of a cluster: with $D_l$ denoting the luminosity distance.
+ By comparing these two expressions we see that. iu contrast to the SZ flux density. the Xrav flux suffers cosmological surface brightuess diuuuiug. represcuted by the extra factors of(1|2) in the denominator of Eq. CU) ," By comparing these two expressions we see that, in contrast to the SZ flux density, the X–ray flux suffers cosmological surface brightness dimming, represented by the extra factors of$(1+z)$ in the denominator of Eq. \ref{eq:fx}) )"
+which couvert the distance to the Inninositv distance., which convert the angular--size distance to the luminosity distance.
+ Besides this wellknown ditfercuce. which tells us that the SZ effect is the more efficicut way to find Lighredshift clusters. we note that the Nray cussion depends ou the eas deusitv in addition to the hot eas mass fraction and temperature.," Besides this well–known difference, which tells us that the SZ effect is the more efficient way to find high–redshift clusters, we note that the X–ray emission depends on the gas density in addition to the hot gas mass fraction and temperature."
+ This is uufortunate. because it imeaus that the Nrav flux from a cluster depends ou the core radius and profile of the intracluster medium (ICAL) two quantities which are poorly. if at all. understood from the theoretical point of view.," This is unfortunate, because it means that the X��ray flux from a cluster depends on the core radius and profile of the intracluster medium (ICM) – two quantities which are poorly, if at all, understood from the theoretical point of view."
+" The SZ dux density presents the iniportaut advantage that it dependstomperatire, aud not on the ICMUSs distribution."," The SZ flux density presents the important advantage that it depends, and not on the ICM's distribution."
+ It is also true that the temperature which appears in the expression for the SZ flux deusitv is a simpler quantity than the Nray iieasured temperature: it is theweighted energy. of tle gas particles instead of. as iu the case of Xrave. the cussionweielted eas temperature.," It is also true that the temperature which appears in the expression for the SZ flux density is a simpler quantity than the X–ray measured temperature: it is the energy of the gas particles instead of, as in the case of X–rays, the emission–weighted gas temperature."
+ This SZ temperature is a quantity which should be all the more closely related to the virial mass of a cluster than even the X.ray temperature. and less affected by any teniperature structure in the cluster.," This SZ temperature is a quantity which should be all the more closely related to the virial mass of a cluster than even the X–ray temperature, and less affected by any temperature structure in the cluster."
+ Now the game is clear: with Eq. (3)), 	Now the game is clear: with Eq. \ref{eq:SZflux}) )
+ we may convert the mass function iuto a distribution of clusters in SZ fiux deusity and redshift (the quantitative relation for Sy. can be found in Barbosa et al. 19963)., we may convert the mass function into a distribution of clusters in SZ flux density and redshift (the quantitative relation for $S_{sz}$ can be found in Barbosa et al. \cite{SZcounts}) ).
+ The redshift distribution of clusters aud the total source counts are thencasily caleulable (I&orolvov ct al. 1986:, The redshift distribution of clusters and the total source counts are theneasily calculable (Korolyov et al. \cite{koro};
+ Alarkevitch et al. 1991: , Markevitch et al. \cite{mark}; ;
+Bartlett Silk 199lb: Barbosa et al. 1996: , Bartlett Silk \cite{bs94}; ; Barbosa et al. \cite{SZcounts}; ;
+Eke ct al., Eke et al.
+ol the bright EUV source HZ 43.,of the bright EUV source HZ 43.
+ The reduced. count rate and large flietuation during this observation (Figure 5¢) is evidently due to this unlortunate. placement. and should be disregarded.," The reduced count rate and large fluctuation during this observation (Figure 5c) is evidently due to this unfortunate placement, and should be disregarded."
+ This QSO was discovered (Reichertetal.1982). as a serendipitous X-ray source in an observation of the RS CVn tvpe star DIL CVn., This QSO was discovered \citep{re82} as a serendipitous X-ray source in an observation of the RS CVn type star BH CVn.
+ It was then detected in an ppointing αἱ (he same star (Christianetal.1999)., It was then detected in an pointing at the same star \citep{ch99}.
+. The small Galactic column in this direction. Vy)=9.2xLOMg >7. undoubtedly contributes to its detectability.," The small Galactic column in this direction, $N_{\rm H} = 9.2 \times 10^{19}$ $^{-2}$, undoubtedly contributes to its detectability."
+ In (his section. we attempt to quantify the statistical significance of (he periodic signals seen in the perioclograms of4711.. Ton S180. and 111 0419577 in Figures 7- 9..," In this section, we attempt to quantify the statistical significance of the periodic signals seen in the periodograms of, Ton S180, and 1H 0419–577 in Figures \ref{ps1}- \ref{ps3}."
+ Perhaps the closest approximation (o a formally correct analvsis of significance is (hat of Timmer&Konig(1995).. as implemented by Benllochetal.(2001).!.," Perhaps the closest approximation to a formally correct analysis of significance is that of \cite{tk95}, as implemented by \cite{bwe01}
+."
+. Timmer(1995) correctly take into account randomness in both Fourier amplitudes and phases in simulating the power spectra of flickering [P(F)xf. 4) or random-walk [P(F)xf. 7] light curves., \cite{tk95} correctly take into account randomness in both Fourier amplitudes and phases in simulating the power spectra of flickering $P(F) \propto f^{-1}$ ] or random-walk $P(F) \propto f^{-2}$ ] light curves.
+" These processes are good descriptions of many AGN power spectra which. in the range of frequencies sampled here. (wpically follow the form PCF)xf""D where 1<a«2."," These processes are good descriptions of many AGN power spectra which, in the range of frequencies sampled here, typically follow the form $P(F) \propto f^{-\alpha}$ where $1 < \alpha < 2$."
+" We apply this method to estimate the statistical significance of tentative periodic signals seen in the periodograms of4711.. Ton S180. and LI 0419577. which are the three Sevlerts that oobserved Lor at least 20 clays each in continuous or nearly continuous stretches,"," We apply this method to estimate the statistical significance of tentative periodic signals seen in the periodograms of, Ton S180, and 1H 0419–577, which are the three Seyferts that observed for at least 20 days each in continuous or nearly continuous stretches."
+ The approach involves simulating may red-noise light curves that match the real data in sampling time. count rate. and variance. and comparing their periodograms to the periodogranms of the real data.," The approach involves simulating many red-noise light curves that match the real data in sampling time, count rate, and variance, and comparing their periodograms to the periodograms of the real data."
+ For a discrete peak in an observed periodogram to be significant. il must have greater power (han nearly all of the simulations at that [reeuency.," For a discrete peak in an observed periodogram to be significant, it must have greater power than nearly all of the simulations at that frequency."
+ In order to determine what red-noise power-law index a matches a particular data set. we followed a procedure similar to that used bv Leishlv.(19994). and Doneetal.(1992).. and essentially the same as Uttlev.Mellardy.&Papacakis(2002).," In order to determine what red-noise power-law index $\alpha$ matches a particular data set, we followed a procedure similar to that used by \cite{l99a}
+ and \cite{do92}, and essentially the same as \cite{ump02}."
+. We simulated 100 light curves for each of a range of values of a and standard deviation σ of the light curve that, We simulated 100 light curves for each of a range of values of $\alpha$ and standard deviation $\sigma$ of the light curve that
+"disparity becomes very large with k,>>K,, where we use ki to distinguish the parallel component of the turbulence wave packet wavenumber ki from the parallel wavenumber of the growing wave kj.","disparity becomes very large with $k_\bot \gg k^t_\|$, where we use $k^t_\|$ to distinguish the parallel component of the turbulence wave packet wavenumber $k^t_\|$ from the parallel wavenumber of the growing wave $k_\|$."
+" In MHD turbulence, wave packet cascades after it travels a distance of e~LIS2>1/k,."," In MHD turbulence, wave packet cascades after it travels a distance of $1/k^t_\| \sim L^{1/3}k_\bot^{-2/3} \gg 1/k_\bot$."
+"Onthe other hand, the instability grows1/ky fastest for the most parallel wave (k~ kj) allowed in a turbulence medium with their wave numbers satisfying Κι/kj~6B/B(ΚιL)-13."," On the other hand, the instability grows fastest for the most parallel wave $k\sim k_\|$ ) allowed in a turbulence medium with their wave numbers satisfying $k_\bot/k_\| \sim \delta B/B\sim (k_\bot L)^{-1/3}$."
+" Because of the scale disparity, kj>k,>Ki the nonlinear damping rate in MHD turbulence is less than the wave frequency ΚΥΑ. and it is given by (Farmer Goldreich 2004; Yan Lazarian 2004; Beresnyak Lazarian 2008) where Lis the injection scale of background turbulence, and the k is set by the resonance condition k~kj1/r,."," Because of the scale disparity, $k_\| > k_\bot \gg k^t_\|$, the nonlinear damping rate in MHD turbulence is less than the wave frequency $k_\| v_A$, and it is given by (Farmer Goldreich 2004; Yan Lazarian 2004; Beresnyak Lazarian 2008) where Lis the injection scale of background turbulence, and the $k$ is set by the resonance condition $k \sim k_\| \sim 1/r_L$."
+" Inserting Eqs.(6,,7)) into Eq.(1)) and adopting UdW/dx+U?W/D in the case of efficient wave amplification, one gets There are various models for the diffusive shock acceleration."," Inserting \ref{growth}, \ref{damping}) ) into \ref{wave}) ) and adopting $U\partial W/\partial x \approx U^2 W/D$ in the case of efficient wave amplification, one gets There are various models for the diffusive shock acceleration."
+ We consider here the escape-limited acceleration., We consider here the escape-limited acceleration.
+" In this model, particles are confined in the region near the shock where turbulence is generated."," In this model, particles are confined in the region near the shock where turbulence is generated."
+" Once they propagate far upstream at a distance / from the shock front, where the self-generated turbulence by CRs fades the and the acceleration ceases."," Once they propagate far upstream at a distance $l$ from the shock front, where the self-generated turbulence by CRs fades away, the particles escape and the acceleration ceases."
+" The away,characteristic particleslength that escapeparticles penetrate into the upstream is D(p)/U.", The characteristic length that particles penetrate into the upstream is $D(p)/U$.
+ The maximum momentum is reached when D(p)/Uc1/4)., The maximum momentum is reached when $D(p)/U\simeq l/4$.
+". Assuming {οςR,,, then the maximum momentum of particles accelerated during the Sedov phase is determined by the condition where K<1 is a numerical factor, see table 2."," Assuming $l\propto R_{sh}$, then the maximum momentum of particles accelerated during the Sedov phase is determined by the condition where $\kappa<1$ is a numerical factor, see table 2."
+" From equations (3,,9)), we get Insert Eq.(10)) into Eq.(8)), we get forA<1° where is approximated by (Pmax/mc)*/a."," From equations \ref{crdiff}, \ref{Dshock}) ), we get Insert \ref{gmax}) ) into \ref{wave2}) ), we get for $A<1$ where $\phi(p_{max})$ is approximated by $(p_{max}/mc)^a/a$."
+" In the limit of 0(Pmax)low shock velocity, we get At the Sedov phase (f>tyeg=250 yr). where Es;= and Uy=(Es,/(noU3)) are the total energy ofEsy/10°!erg explosion and the initialU;/10?cm/s shock velocity, the evolution of shock radius and speed are governed by From Eqs.(13,,14)), we see that Pinayοςt9, and 6=18/5."," In the limit of low shock velocity, we get At the Sedov phase $t>t_{sed}\equiv 250(E_{51}/(n_0U_9^5))^{1/3}$ yr), where $E_{51}=E_{SN}/10^{51}$ erg and $U_9=U_i/10^9$ cm/s are the total energy of explosion and the initial shock velocity, the evolution of shock radius and speed are governed by From \ref{gmax_solution}, \ref{Sedov}) ), we see that $p_{max} \propto t^{-\delta}$, and $\delta=18/5$."
+" In Fig.1,, we plot Pmax/(mc) vs. the time t since supernova explosion."," In \ref{Emax},, we plot $p_{max}/(mc)$ vs. the time t since supernova explosion."
+ The solid line represents the results from Eqs.(11))., The solid line represents the results from \ref{gmax_general}) ).
+" As we see, at earlier epoch when advection and streaming instability are both important, the evolution of Ping, does not follow a power law."," As we see, at earlier epoch when advection and streaming instability are both important, the evolution of $p_{max}$ does not follow a power law."
+" For comparison, we also put a power law evolution in the same figure as depicted by Eq.(13)) (dashed line)."," For comparison, we also put a power law evolution in the same figure as depicted by \ref{gmax_solution}) ) (dashed line)."
+ Our result is also larger than that obtained by Ptuskin Zirakashvili (2005) since the wave dissipation rate is overestimated in their treatment., Our result is also larger than that obtained by Ptuskin Zirakashvili (2005) since the wave dissipation rate is overestimated in their treatment.
+" At a given time, with p>pmax can escape the shock front."," At a given time, only particles with $p>p_{max}$ can escape the shock front."
+" Adopting onlythe particlessimplified approximation, at a given energy E—cp, there is a one to one correspondence between the CR momentum and the CR escape time ο (or radius Resp(14- K)Rs;) in spite of the fact that the acceleration is continuous the shock expansion since reaches p."," Adopting the simplified approximation, at a given energy $E = cp$, there is a one to one correspondence between the CR momentum and the CR escape time $t_{esp}$ (or radius $R_{esp}=(1+\kappa)R_{sh}$ ) in spite of the fact that the acceleration is continuous during the shock expansion since $p_{max}(t)$ reaches $p$ ."
+" If the duringmaximum momentum has a power law Pmax(t)dependence on t, one can easily gets tespοςp 9.Since we consider the later stage of SNR acceleration, the shock radius cannot be neglected and the point source"," If the maximum momentum has a power law dependence on t, one can easily gets $t_{esp}\propto p^{-1/\delta}$ .Since we consider the later stage of SNR acceleration, the shock radius cannot be neglected and the point source"
+for their AGN power compared to the optically selected PG QSOs.,for their AGN power compared to the optically selected PG QSOs.
+ An interesting trend in Table 1 is the high fraction of the 2MASS QSOs with detection at GGHz: 6/10 are radio sources in either the NVSS of the the FIRST surveys., An interesting trend in Table 1 is the high fraction of the 2MASS QSOs with detection at GHz: 6/10 are radio sources in either the NVSS of the the FIRST surveys.
+" This is unexpected because the sources were selected independent of their radio properties, unlike previous efforts to identify reddened QSOs, which started from radio selected samples to identify objects with very red optical/near-IR colours (e.g.2222?).."," This is unexpected because the sources were selected independent of their radio properties, unlike previous efforts to identify reddened QSOs, which started from radio selected samples to identify objects with very red optical/near-IR colours \citep[e.g.][]{Gregg2002, White2003, Glikman2004, Glikman2007,
+ Urrutia2008a}."
+ Next we explore whether the radio detection rate of 2MASS QSOs is different from optically selected ones., Next we explore whether the radio detection rate of 2MASS QSOs is different from optically selected ones.
+" For this exercise, we compare against the SDSS QSOs to take advantage of the large sample size and the well defined selection criteria."," For this exercise, we compare against the SDSS QSOs to take advantage of the large sample size and the well defined selection criteria."
+" We use in particular the SDSS-DR3 quasar catalog of ?,, which was created by inspecting all spectra that were either targeted as quasar candidates, or classified as a quasar by the SDSS spectroscopic pipelines."," We use in particular the SDSS-DR3 quasar catalog of \cite{Schneider2005}, which was created by inspecting all spectra that were either targeted as quasar candidates, or classified as a quasar by the SDSS spectroscopic pipelines."
+" This catalogue includes objects from all categories of spectroscopic target selection in the SDSS, not just those selected as QSO candidates."," This catalogue includes objects from all categories of spectroscopic target selection in the SDSS, not just those selected as QSO candidates."
+" For our purposes we use only targeted SDSS QSOs in the ? catalogue, i.e. sources that belong to the the main (or low-redshift) and the high-redshift QSO samples of the SDSS."," For our purposes we use only targeted SDSS QSOs in the \cite{Schneider2005} catalogue, i.e. sources that belong to the the main (or low-redshift) and the high-redshift QSO samples of the SDSS."
+ These samples include QSOs with 15<7« 19mmag that were selected based on the optical colour criteria of ?.., These samples include QSOs with $15<i<19$ mag that were selected based on the optical colour criteria of \cite{Richards2002}.
+ We further apply acut kK<14.5 mmag to match the magnitude limit used to select 2MASS QSOs., We further apply a cut $K<14.5$ mag to match the magnitude limit used to select 2MASS QSOs.
+ There are 638 such sources in the ? catalogue., There are 638 such sources in the \cite{Schneider2005} catalogue.
+" The K—14.5 mmag magnitude cut does not introduce incompleteness in the sample, as it corresponds to {<17.5 mmag for the average (i—K) colour of optically selected QSOs, which takes values in the range 2-3, depending on redshift."," The $K=14.5$ mag magnitude cut does not introduce incompleteness in the sample, as it corresponds to $i \la
+17.5$ mag for the average $(i-K)$ colour of optically selected QSOs, which takes values in the range 2–3, depending on redshift."
+" Similarly, the bright magnitude limit 2> 15mmag for follow up spectroscopy of the SDSS QSO candidates corresponds to K=12.5 mmag."," Similarly, the bright magnitude limit $i>15$ mag for follow up spectroscopy of the SDSS QSO candidates corresponds to $K \ga
+12.5$ mag."
+ QSOs brighter than that may be missed from the SDSS sample., QSOs brighter than that may be missed from the SDSS sample.
+" The surface density of such bright QSOs is likely to be small however, and it is not expected to introduce any significant bias in the SDSS optically selected QSO sample."," The surface density of such bright QSOs is likely to be small however, and it is not expected to introduce any significant bias in the SDSS optically selected QSO sample."
+ In order to estimate the radio detection rate of the sample we use the FIRST radio survey., In order to estimate the radio detection rate of the sample we use the FIRST radio survey.
+ There are 582 SDSS QSOs with K«14.5 mmag that fulfill the ? criteria and overlap with the area covered by the FIRST survey., There are 582 SDSS QSOs with $K<14.5$ mag that fulfill the \cite{Richards2002} criteria and overlap with the area covered by the FIRST survey.
+" A total of 172 of them, i.e. =30 per cent, are detected at GGHz to the FIRST limit."," A total of 172 of them, i.e. $\approx30$ per cent, are detected at GHz to the FIRST limit."
+" This fraction should be compared with the detection rate of 2MASS QSOs, i.e. 60 per cent."," This fraction should be compared with the detection rate of 2MASS QSOs, i.e. 60 per cent."
+ Using binomial statistics we estimate a 4 per cent probability of at least 6 radio detections in 10 trials given the success rate of 30 per cent of the optically selected QSO sample., Using binomial statistics we estimate a 4 per cent probability of at least 6 radio detections in 10 trials given the success rate of 30 per cent of the optically selected QSO sample.
+" The radio detection rate of 2MASS QSOs is higher than that of optically selected ones, albeit at the 96 per cent level"," The radio detection rate of 2MASS QSOs is higher than that of optically selected ones, albeit at the 96 per cent level"
+ =(ALv», =(M_1.
+", The chirp mass of a NS-NS binary. with both neutron stars of mass 1.4AL... ismsun."," The chirp mass of a NS-NS binary, with both neutron stars of mass $1.4\msun$, is."
+. A birth rate of CGIEM implics a rate of tout to 200MAMpe (Phniunev 1991)., A birth rate of GEM implies a rate of 3 $^{-1}$ out to Mpc (Phinney 1991).
+ Kip Thorne informs us that LICO’s first loue eravitationalwave search. in 20023 as discussed for binary neutron stars ds expected to see binaries with Aa=1.2M. out to AAIpec.," Kip Thorne informs us that LIGO's first long gravitational-wave search in $-$ 2003 as discussed for binary neutron stars is expected to see binaries with $M_{{\rm
+chirp}}=1.2\msun$ out to Mpc."
+ The chirp mass corresponding to the Bethe Brown (1995) LDII-NSbinary with masses 2.1M. and 1.1AY... respectively.isAZ.," The chirp mass corresponding to the Bethe Brown (1998) LBH-NS binary with masses $2.4\msun$ and $1.4\msun$ , respectively, is."
+.. Iucludiug a ~30% increase in the rate to allow for ligh-mass black-hole ΠΟΠ)ο ierecrs (Bethe Brown 1999) eives a 26 times higher rate than Phiunevs estimate for NS-NS mereers (10.7 + in the Galaxy)., Including a $\sim 30\%$ increase in the rate to allow for high-mass black-hole (HBH)-NS mergers (Bethe Brown 1999) gives a 26 times higher rate than Phinney's estimate for NS-NS mergers $10^{-5}$ $^{-1}$ in the Galaxy).
+" These factors are calculated from the signal to noise ratio. which goes as MASS, and then cubing it to obtain the volume of detectal(lits. which is therefore proportional to MAL."," These factors are calculated from the signal to noise ratio, which goes as $M_{{\rm chirp}}^{5/6}$, and then cubing it to obtain the volume of detectability, which is therefore proportional to $M_{{\rm chirp}}^{5/2}$."
+ We then predict a rate of 3«(21/200)?«26=0.09 1., We then predict a rate of $3\times(21/200)^3\times26=0.09$ $^{-1}$.
+ This rate is slim for 2003., This rate is slim for 2003.
+ The chhanced LIGO interferometer planned to begin in 2001 should reach out bevoud MMpc for »200)?hirp—1.2AL... increasing the detection rate to 3<n(15t.26233 ft. and IIBUENS mereers used estimates should be considered a lower linut (Sect. mthe 53.2)).," The enhanced LIGO interferometer planned to begin in 2004 should reach out beyond Mpc for $M_{{\rm chirp}}=1.2\msun$, increasing the detection rate to $3\times(150/200)^3\times 26 =33$ $^{-1}$, and HBH-NS mergers used in these estimates should be considered a lower limit (Sect. \ref{pop.WR}) )."
+ We therefore fiud that inclusion of black holes iu estimates for LIGO predict that we will ee more merecrs per mouth than NS-NS mergers per vear., We therefore find that inclusion of black holes in the estimates for LIGO predict that we will see more mergers per month than NS-NS mergers per year.
+ Because eamuna-ray bursts lave a median redshitt of 1.52 WWijers et 11998). and the supernova rate at that redshift was 1020 times higher Έλα now. the eamunaray burst rate as observed ds ligher than oue expects uxiug the above rates.," Because gamma-ray bursts have a median redshift of 1.5–2 Wijers et 1998), and the supernova rate at that redshift was 10–20 times higher than now, the gamma-ray burst rate as observed is higher than one expects using the above rates."
+ Eowever. for ease of comparison with evolutionary scenarios we shall use the GRB rate at the preseut time (redshift 0) of about GGEN. (," However, for ease of comparison with evolutionary scenarios we shall use the GRB rate at the present time (redshift 0) of about GEM. ("
+Wijers et (1998) found a factor 3 lower rate. but had slightly uuderestinated it because thev overestimated the mean CRB redshift: see Fever. Woosley. Tlartinann (1999) for more exteusive discussions of the redshift dependence}.,"Wijers et (1998) found a factor 3 lower rate, but had slightly underestimated it because they overestimated the mean GRB redshift; see Fryer, Woosley, Hartmann (1999) for more extensive discussions of the redshift dependence)."
+ An important uncertainty is the beanuue of eamuna-ray bursts: the eamuna ravs may only be eumütted i narrow cones around the spin axis of the black hole. aud therefore most CRBs may not be secu by us.," An important uncertainty is the beaming of gamma-ray bursts: the gamma rays may only be emitted in narrow cones around the spin axis of the black hole, and therefore most GRBs may not be seen by us."
+ An upper liit to the ratio of uudoetected to detected GRB is 600 (Mésszirros. Rees. Wijers 1999). so au mye laut to the total required formation rate would be GEM.," An upper limit to the ratio of undetected to detected GRB is 600 (Mésszárros, Rees, Wijers 1999), so an upper limit to the total required formation rate would be GEM."
+ We may have secu beaming of about that factor or a bit less in 9990125 (I&ulkarni et al., We may have seen beaming of about that factor or a bit less in 990123 (Kulkarni et al.
+ 1999). but other bursts 9970228. 970508) show no evidence of beaming iu the afterglows (which may not exclude beamine of their eamuna rays).," 1999), but other bursts 970228, 970508) show no evidence of beaming in the afterglows (which may not exclude beaming of their gamma rays)."
+ At present. therefore. ay progenitor with a formation rate of LOGCGEAL or more should be considered consistent with the observed. CRD rate.," At present, therefore, any progenitor with a formation rate of GEM or more should be considered consistent with the observed GRB rate."
+ We have shown that rapidly rotating black holes are an attractive power source for @auunarav bursts, We have shown that rapidly rotating black holes are an attractive power source for gamma-ray bursts.
+ Via the Blandford-Zuajels mechanisia (1977) they can supply suffüiiieut energv at a lieh rate., Via the Blandford-Znajek mechanism (1977) they can supply sufficient energy at a high rate.
+ They also occur often enough to explain the observed CRB rate. even if the eanuni-ray Cluission of a typical GRD is beamed to less of the sky.," They also occur often enough to explain the observed GRB rate, even if the gamma-ray emission of a typical GRB is beamed to less of the sky."
+ Because of the requirement of rapid spin. the direct collapse of a stellar core to a black hole is a less likely candidate for making GRD (at least via the BZ effect).," Because of the requirement of rapid spin, the direct collapse of a stellar core to a black hole is a less likely candidate for making GRB (at least via the BZ effect)."
+ Therefore. merecrs are much more attractive. which nuplies a natural connection between CRBs aud strong sources of eravity waves.," Therefore, mergers are much more attractive, which implies a natural connection between GRBs and strong sources of gravity waves."
+ With advanced LIGO. the detection rate of merecrs is predicted to become large enough that direct verification of eveuts that produce both eravity wave aud gauumui-rav signals will become feasible. and will directly coustrain CRB beaming.," With advanced LIGO, the detection rate of mergers is predicted to become large enough that direct verification of events that produce both gravity wave and gamma-ray signals will become feasible, and will directly constrain GRB beaming."
+ We have used the population svuthesis calculations of Bethe Brown (1998. 1999) to estimate the LIGO detection rate.," We have used the population synthesis calculations of Bethe Brown (1998, 1999) to estimate the LIGO detection rate."
+ We found it to be dominated by black-hole. neutrou-star merecrs. and to be higher by a factor 26 than previous estimates.," We found it to be dominated by black-hole, neutron-star mergers, and to be higher by a factor 26 than previous estimates."
+ As a result. we conclude that the most energetic phenomena in astrophysics stem from black holes. whose defining characteristic is paradoxically that no radiation can escape frou then.," As a result, we conclude that the most energetic phenomena in astrophysics stem from black holes, whose defining characteristic is paradoxically that no radiation can escape from them."
+ We would like to thank Roger Blaudford. Chris Fiver. Sterl Phinney. Sinou Portegies Zwart. Kip Thorue aud Stau Woosley for useful sugeestionus and advice.," We would like to thank Roger Blandford, Chris Fryer, Sterl Phinney, Simon Portegies Zwart, Kip Thorne and Stan Woosley for useful suggestions and advice."
+ This work was partially supported bv the U.S. Departineut ofEncrev under Caant No., This work was partially supported by the U.S. Department ofEnergy under Grant No.
+ DEΕΣssER10358., DE–FG02–88ER40388.
+ TIS is also supported partly by NOSEF 955-0200-001-2., HKL is also supported partly by KOSEF 985-0200-001-2.
+The Wall term: does not affect the time evolution of B. siuce V«(JB)x.WSVB? vanishes.,"The Hall term does not affect the time evolution of $\bb$, since $\bnabla \times (\bj \times \bb) \propto -\bnabla \times \bnabla B^2$ vanishes."
+ Thus. ambipolar diffision acts asresistivity.," Thus, ambipolar diffusion acts as."
+ We define the ambipolar ciffusion coefficient as where Be; micans D ii gauss., We define the ambipolar diffusion coefficient as where $B_G$ means $B$ in gauss.
+ For our model of the protoplauctary disk. we take thenüunnuuniduass solar nebula of IHawvashi(1981).," For our model of the protoplanetary disk, we take theminimum-mass solar nebula of \citet{hay81}."
+" Tt is characterized by the surface mass density “=Sy(R/BAU)77 with Xy=23«109 ο 7. where Rois the radial distance from the central star. and the temperature profile given by T,,=1.6sJUCR/AAU)3,"," It is characterized by the surface mass density $\Sigma = \Sigma_0 \, (R/3 \AU)^{-3/2}$ with $\Sigma_0 = 3.3 \times 10^2$ g $^{-2}$, where $R$ is the radial distance from the central star, and the temperature profile given by $T_g= 1.6 \times 10^2 \, (R/3 \AU)^{-1/2}$."
+ We expect Xy to decrease in time as the nebular eas clears out iu the late stage of protoplanctary disk evolution., We expect $\Sigma_0$ to decrease in time as the nebular gas clears out in the late stage of protoplanetary disk evolution.
+ Hence we use the form Sy=3:3:&lO? ο 7. where fe is an arbitrary parameter less than or equal to unity.," Hence we use the form $\Sigma_0 = 3.3 \times 10^2 f_{\Sigma}$ g $^{-2}$, where $f_{\Sigma}$ is an arbitrary parameter less than or equal to unity."
+" If we take the eas to be isothermal. the vertical density profile in hwdrostatie equilibrium has a Caussian form (Ixeuvon(9) where iy=MAY2a10,,h)l.l«1015Fe(R/3AU)Yt ? is the umber deusity of gas in the disk midplane. : is the vertical distance measured from the midplane. aud /—CsQO is the scale height."," If we take the gas to be isothermal, the vertical density profile in hydrostatic equilibrium has a Gaussian form \citep[][hereafter CG97]{ken87,chi97}
+ where $n_0 = \Sigma / (\sqrt{2\pi} \, m_n h) = 1.4 \times 10^{13} \, f_{\Sigma} \, (R/3 \AU)^{-11/4}$ $^{-3}$ is the number density of gas in the disk midplane, $z$ is the vertical distance measured from the midplane, and $h \equiv c_s/\Omega$ is the scale height."
+" Here. as usual ¢=(KT,nq)H7 ds the souud speed and © is the IKepleriau rotation rate."," Here, as usual, $c_s = (k T_g/m_n)^{1/2}$ is the sound speed and $\Omega$ is the Keplerian rotation rate."
+ With these values. we cau roughly estimate the Óuportance of ambipolar diffusion by taking the ratio of the two coefficients iu. Equ. (CL):," With these values, we can roughly estimate the importance of ambipolar diffusion by taking the ratio of the two coefficients in Eqn. \ref{ind2}) ):"
+ We note that (1) with a magnetic fold sufBcieutlv aipli&ed (up to 1 €). ambipolar diffusion is at least as dnportant as Olunic dissipation and (2) the ratio las strong spatial dependence within the disk.," We note that (1) with a magnetic field sufficiently amplified (up to $\sim$ 1 G), ambipolar diffusion is at least as important as Ohmic dissipation and (2) the ratio has strong spatial dependence within the disk."
+ Remuaut uaenetiuu in meteorites suggests that B2: 1 Ce was vpical in the solar nebula., Remnant magnetism in meteorites suggests that $B \approx$ 1 G was typical in the solar nebula.
+ Tf an efficient. dynamo is in action. the ambipolar diffusion termi would eeucrallv be important iu the outer (Rz 1 AU) region of the disk.," If an efficient dynamo is in action, the ambipolar diffusion term would generally be important in the outer $R \gtrsim$ 1 AU) region of the disk."
+ In he inner region. for a simular value of D aud fe=1l. it is nouneelieible only a few scale heights above aud below he midplane. where gas density is nach lower.," In the inner region, for a similar value of $B$ and $f_{\Sigma}=1$, it is nonnegligible only a few scale heights above and below the midplane, where gas density is much lower."
+ It ecucrally )ocomies Lacreasinely important as fe decreases., It generally becomes increasingly important as $f_{\Sigma}$ decreases.
+ Compared to other astroplivsical svstenis. protoplanctary disks are unuusuallv cold aud dense and are therefore associated with very low ionization fractions.," Compared to other astrophysical systems, protoplanetary disks are unusually cold and dense and are therefore associated with very low ionization fractions."
+ louizatiou fraction is a key parameter in determining whether a disk will be turbulent or not. since the MRI is confined to regious where the ionization fraction is high choueh to cusure eood coupling between ious aud ucutrals.," Ionization fraction is a key parameter in determining whether a disk will be turbulent or not, since the MRI is confined to regions where the ionization fraction is high enough to ensure good coupling between ions and neutrals."
+" This occurs when the ion-neutral collision rate (55;,) 1s sjeuificautlv ereater (bv 3100 at least) than the rotation rate (OQ) (Blacs&Balbus1991:MacLowctal.1995:Hawley&Stone1998)."," This occurs when the ion-neutral collision rate $\nu_{in}$ ) is significantly greater (by 3–100 at least) than the rotation rate $\Omega$ ) \citep{bla94,mac95,haw98}."
+. The critical ionization fraction is calculated to he (e.g...Balbus&Hawley1998).," The critical ionization fraction is calculated to be \citep[e.g.,][]{bal98}."
+. Iuterestiugly. this value is ouly slightly higher than the values of LO12 10i? estimated in the bulk of protoplauetary disks (Ctaunnie1996:Classeold.Najita.&Igea1997:Sanoetal. 2000).," Interestingly, this value is only slightly higher than the values of $10^{-12}$ $10^{-13}$ estimated in the bulk of protoplanetary disks \citep{gam96,gla97,san00}."
+. Many nmnnerieal simulations have demonstrated that AIRLiuduced velocity shears amplify the maguetic field (see.ec.Hawley&Balbus1991:Draudeubursctal.1995:Tlawlev.Cammue.&Dalbus 1995).," Many numerical simulations have demonstrated that MRI-induced velocity shears amplify the magnetic field \citep[see, e.g.,][]{haw91,brae95,haw95}."
+. Tn our model. perturbations in the magnetic feld are also assmued to be generated by the MRI.," In our model, perturbations in the magnetic field are also assumed to be generated by the MRI."
+ Considering the low values of Y; estimated in protoplauctary disks. it is iuitiallv unclear whether this assumption is valid or not.," Considering the low values of $\chi_i$ estimated in protoplanetary disks, it is initially unclear whether this assumption is valid or not."
+ Therefore we ask: which resion of the solar nebula. if any. is prone to the MRI? (," Therefore we ask: which region of the solar nebula, if any, is prone to the MRI? ("
+Previous numerical simulations assumed ionization fraction X;2νε),Previous numerical simulations assumed ionization fraction $\chi_i \gg \chi_{crit}$ .)
+ This question bas beeu investigated by several authors (Uiuehavashi1983:Sanoetal2000:Fromane.Terquem.&Balbus 2002).," This question has been investigated by several authors \citep{ume83,san00,fro02}."
+.. Their results are as follows., Their results are as follows.
+" The ionization fraction 1s sutiicienutly lieh iu the imuermost region of the disk up to a few tenths of an AU. where the ambicut gas teniperature is high cuough (T,=10° TS) for collisional ionization to be effective (whereitdominatesovercosmicτανaudracdioac-tiveclementionizatious:Uiiebavashi 1983)."," The ionization fraction is sufficiently high in the innermost region of the disk up to a few tenths of an AU, where the ambient gas temperature is high enough $T_g \gtrsim 10^3$ K) for collisional ionization to be effective \citep[where it dominates over cosmic ray and radioactive element ionizations;][]{ume83}."
+". Also. cosmic ravs and stellar X-ray radiation iouize the surface lavers (vith thickness given by the stopping depth of X,z10? efau?) at R<10 AU aud the cutive disk at R>10 AU. where the surface density of the disk falls below στα."," Also, cosmic rays and stellar X-ray radiation ionize the surface layers (with thickness given by the stopping depth of $\Sigma_a \approx 10^2$ $^2$ ) at $R \lesssim 10$ AU and the entire disk at $R \gtrsim 10$ AU, where the surface density of the disk falls below $\Sigma_a$ ."
+ These. then. are the regions where the disk is turbulent. rence effectively transterring augular momentum and also actively accreting.," These, then, are the regions where the disk is turbulent, hence effectively transferring angular momentum and also actively accreting."
+ On the other laud. the region near the disk amidplane. where terrestrial plaucts aud the asteroid olt now lie. is generally thought to contain no siguificaut ιοπιο source and be quiesceut. ie. magnoeticallv dead (radioactiveheatingistooineficicuttomaintainX;> Venti1996).," On the other hand, the region near the disk midplane, where terrestrial planets and the asteroid belt now lie, is generally thought to contain no significant heating source and be quiescent, i.e. magnetically dead \citep[radioactive heating is too inefficient to maintain $\chi_i > \chi_{crit}$ ;."
+ These considerations led to the avered accretion model of Caamuiue (1996:seealsogold.Najita.&Igea 1997)., These considerations led to the layered accretion model of Gammie \citep[1996; see also][]{gla97}.
+. There are two leugth scales. other than the scale height h. of relevance to this cliscussion.," There are two length scales, other than the scale height $h$, of relevance to this discussion."
+ The wavelength of perturbation most unstable to the MBRI is only slightly larger than the critical waveleugth at which the instability sets in (BIT9L) ancl is given by (BII91)., The wavelength of perturbation most unstable to the MRI is only slightly larger than the critical wavelength at which the instability sets in (BH91) and is given by (BH91).
+ Other processes affect the erowth of perturbations: iu dense parts of the disk. collisious with ueutrals binder the charged compoucuts from drifting with the maenetic field.," Other processes affect the growth of perturbations: in dense parts of the disk, collisions with neutrals hinder the charged components from drifting with the magnetic field."
+" In this case. Olunic dissipation acts to stabilize wavelengths shorter than the resistive diffusion scale My= 224/04. where eq=Bf(1xp,)!7 is the Alfven speed."," In this case, Ohmic dissipation acts to stabilize wavelengths shorter than the resistive diffusion scale $\lambda_{\Omega} \equiv 2 \pi \eta /v_A$ , where $v_A=B/(4\pi \rho_n)^{1/2}$ is the Alfven speed."
+ Iu other words. the critical waveleneth at which MRI sets in isgiven by Αη=ae(Apy.Ao) ," In other words, the critical wavelength at which MRI sets in isgiven by $\lambda_{crit} = max \, (\lambda_{BH},\lambda_{\Omega})$ "
+"Note that in quiescence, the NIR/optical of iis likely dominated by an optically thin disc plus a much fainter mass donor star (?)..","Note that in quiescence, the NIR/optical of is likely dominated by an optically thin disc plus a much fainter mass donor star \citep{Shab01}."
+" Finally,. we had sufficient. NIR and optical. data to ascertain. when wwas in the soft state and constructed the corresponding SED around 55303.7 MJD (Fig. 10))."," Finally, we had sufficient NIR and optical data to ascertain when was in the soft state and constructed the corresponding SED around 55303.7 MJD (Fig. \ref{sedfarid}) )."
+" In fitting the SED, we used a technique similar to the one described in ?:: we tested a model consisting of a (??) added to aGAUSSIAN, versus a model consisting of: +GAUSSIAN +COMPTT."," In fitting the SED, we used a technique similar to the one described in \citet[]{Rahoui10}: we tested a model consisting of a \citep{Gierlinski:2008,gier2009} added to a, versus a model consisting of: + +."
+ We fixed kT. to the temperature found in Table 2.., We fixed $kT_{\rm e}$ to the temperature found in Table \ref{tab:para}.
+" We found optical s|opes around 2, which are consistent with thermal radiation."," We found optical slopes around 2, which are consistent with thermal radiation."
+ The reprocessing levels in the disc would need to be very high, The reprocessing levels in the disc would need to be very high
+We use Equation 6 το estimate the 1.4. Gllz Dux density emitted. by cach wind phase. assuming a uniform. spherically-svnimetrie wind: the values are listed in Table 1 alongside the associated mass outflow rates calculated using Equation 4..,"We use Equation \ref{s_nu_WB} to estimate the 1.4 GHz flux density emitted by each wind phase, assuming a uniform, spherically-symmetric wind; the values are listed in Table \ref{wa_properties_results}, alongside the associated mass outflow rates calculated using Equation \ref{mdot_4pi}."
+ For the purposes of this calculation. we assume that po= 1.23. 5=1. Z=L4 and extrapolate qy from. Figure 5 in WKarzas&Latter(1961).," For the purposes of this calculation, we assume that $\mu=1.23$ , $\gamma=1$, $\bar Z=1.4$ and extrapolate $g_{\rm ff}$ from Figure 5 in \citet{karzas1961}."
+.. H0 turns out that. generally. only the phases with the highest mass outflow rates are predicted to produce radio ux densities greater jui or equal to the observed. values: for these. phases. we ise Equation 9. to obtain an upper limit to CV.," It turns out that, generally, only the phases with the highest mass outflow rates are predicted to produce radio flux densities greater than or equal to the observed values; for these phases, we use Equation \ref{cv_omega} to obtain an upper limit to $C_{\rm v} \Omega$."
+ An exception is MCX-6-30-15. where the highest radio —ux is predicted for the non-outfowing phase with the lowest --onisation parameter.," An exception is MCG-6-30-15, where the highest radio flux is predicted for the non-outflowing phase with the lowest ionisation parameter."
+ In this source. the partially-covering ise wind component. (phase 5) could be a feasible source M a large part of the observed. radio flux density. it has v volume filling factor of near unity.," In this source, the partially-covering disc wind component (phase 5) could be a feasible source of a large part of the observed radio flux density, it has a volume filling factor of near unity."
+ LE on the other hand. uns volume filling factor is comparable to that of phase 2. which must be less than 0.02 (see Table 1)). then this is unlikely.," If, on the other hand, its volume filling factor is comparable to that of phase 2, which must be less than 0.02 (see Table \ref{wa_properties_results}) ), then this is unlikely."
+ In the case of NGC 3783. none of the wind. phases can produce a significant fraction of the observed. radio Dux density. although the combined. Hux of all. phases comes close.," In the case of NGC 3783, none of the wind phases can produce a significant fraction of the observed radio flux density, although the combined flux of all phases comes close."
+ The global covering factor 2 of AGN winds is uncertain., The global covering factor $\Omega$ of AGN winds is uncertain.
+ lt is known that ~50% of nearby Sevlert Ls have a wind (Revnolds1997).. implving a covering [actor of 2z. which would. be entirely. possible for an accretion disc wind.," It is known that $\sim$ of nearby Seyfert 1s have a wind \citep{reynolds1997}, implying a covering factor of $2\pi$, which would be entirely possible for an accretion disc wind."
+ {4 seems increasinglv likely. however. that most of the observed absorbing column in the winds is associated with gas at the distance of the torus or NLR. implving that the covering actor cannot be greater than the solid angle of the openings in the torus.," It seems increasingly likely, however, that most of the observed absorbing column in the winds is associated with gas at the distance of the torus or NLR, implying that the covering factor cannot be greater than the solid angle of the openings in the torus."
+ The minimum covering factor will therefore be he product ofthe percentage of AGN that are unobscured (i.c. the fraction of AGN that are type 1: 254) and. the »reentage of type 1. AGN with warm. absorbers. giving QO=1.6.," The minimum covering factor will therefore be the product of the percentage of AGN that are unobscured (i.e. the fraction of AGN that are type 1; ) and the percentage of type 1 AGN with warm absorbers, giving $\Omega=1.6$."
+ For the purposes of our calculations. we use these wo extremes of O to give a range of έν ," For the purposes of our calculations, we use these two extremes of $\Omega$ to give a range of $C_{\rm v}$."
+The overall result is that. for the reasonable range of ©. he upper limits to Cy are of the order of 101o(L5: the igures are listed in Table 1," The overall result is that, for the reasonable range of $\Omega$, the upper limits to $C_{\rm v}$ are of the order of $10^{-4}-0.5$; the figures are listed in Table \ref{wa_properties_results}."
+ We have shown that. at least. for some phases of X-ray absorbing winds. the volume filling factors need to be small: of the order of 10.!5.," We have shown that, at least for some phases of X-ray absorbing winds, the volume filling factors need to be small: of the order of $10^{-4}-0.5$."
+ They may well be lower than this. since we have not taken into account the radio emission [roni the UV-absorbing part of the outflow. and we have assumed that there is no contribution from the host galaxies. the base of any radio jet. or indeed emission from the accretion disc corona (asproposedbyLaor&Behar2008).," They may well be lower than this, since we have not taken into account the radio emission from the UV-absorbing part of the outflow, and we have assumed that there is no contribution from the host galaxies, the base of any radio jet, or indeed emission from the accretion disc corona \citep[as~proposed~by][]{laor2008}."
+. Certainly. in the case of NGC 3783 where we do not predict the X-ray absorbing wind to be a feasible source of the observed. racio emission. the fact that the radio source is slightly extended may imply that emission from the host galaxy is a significant component.," Certainly, in the case of NGC 3783 where we do not predict the X-ray absorbing wind to be a feasible source of the observed radio emission, the fact that the radio source is slightly extended may imply that emission from the host galaxy is a significant component."
+ Any future observations of these sources which resolve jet or host galaxy components would allow us to further constrain the volume filling factors., Any future observations of these sources which resolve jet or host galaxy components would allow us to further constrain the volume filling factors.
+ Η the filling factors of the wind. phases are small. it is interesting to estimate the radio emission from the plasma illing the rest of the space.," If the filling factors of the wind phases are small, it is interesting to estimate the radio emission from the plasma filling the rest of the space."
+ X. possible scenario is that he relatively low-ionisation phases with measurable filling actors are embedded. in a hot medium. outllowing at the same speed. with which they are in pressure equilibrium.," A possible scenario is that the relatively low-ionisation phases with measurable filling factors are embedded in a hot medium, outflowing at the same speed, with which they are in pressure equilibrium."
+ boa his case. assuming. for cxample. that 7=QI for the 100 plasma. we find that the radio Hux densities due to the confining media are at least two orders of magnitude lower han the observed radio Hux densities for cach source.," In this case, assuming, for example, that $T=10^8$ K for the hot plasma, we find that the radio flux densities due to the confining media are at least two orders of magnitude lower than the observed radio flux densities for each source."
+ The radio emission from. raclio-quict AGN can. be variable on timescales of a month or less (Barvainisctal.2005).. which is probably much shorter than the timescale of significant changes to the mass outllow rate in parsec-scale AGN winds.," The radio emission from radio-quiet AGN can be variable on timescales of a month or less \citep{barvainis2005}, which is probably much shorter than the timescale of significant changes to the mass outflow rate in parsec-scale AGN winds."
+ On the other hand. if some fraction of the radio emission originated. [rom a partiallv-covering cise wind. and if the mass outllow rate and. therefore racio Hux varied according to how much of the source the wind was covering. then correlated variations in radio lux ancl absorber partial covering could. provide observational support for the existence of such wind components.," On the other hand, if some fraction of the radio emission originated from a partially-covering disc wind, and if the mass outflow rate and therefore radio flux varied according to how much of the source the wind was covering, then correlated variations in radio flux and absorber partial covering could provide observational support for the existence of such wind components."
+ AJD and ACT acknowledge the support of. respectively. an οΕς Postdoctoral Fellowship and the Itoval Society.," AJB and ACF acknowledge the support of, respectively, an STFC Postdoctoral Fellowship and the Royal Society."
+ We thank L. Miller for providing extra information about the analysis of Milleretal.(2008)., We thank L. Miller for providing extra information about the analysis of \citet{miller2008}.
+. This research has mace use ofthe NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Administration.," This research has made use ofthe NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+55 (νο reported bv Ixórrding et al. (,SS Cyg reported by Körrding et al. (
+2008) originated [rom magnetic activity Chat formed a corona similar to coronae found in magnetically active stars. rather than [from jets.,"2008) originated from magnetic activity that formed a corona similar to coronae found in magnetically active stars, rather than from jets."
+" Soker Vrtilek (2009) based their claim on the results of Laor Behar (2008: hereafter LD2003). who found that when the ratio between radio ancl X-ray fluxes of accretion disks in quasars is as in active stars. L,/L,<10. then most of the radio emission might come from coronae."," Soker Vrtilek (2009) based their claim on the results of Laor Behar (2008; hereafter LB2008), who found that when the ratio between radio and X-ray fluxes of accretion disks in radio-quiet quasars is as in active stars, $L_r/L_x \la 10^{-5}$, then most of the radio emission might come from coronae."
+ Jets might still occur., Jets might still occur.
+ If (he magnetic activity in erupting accreting disks is similar {ο (hat in active stars. then mass ejeclion. is expected.," If the magnetic activity in erupting accreting disks is similar to that in active stars, then mass ejection, is expected."
+ The presence of coronae above accretion disks (e.g.. Galeev οἱ al.," The presence of coronae above accretion disks (e.g., Galeev et al."
+ 1979: Done Osborne 1997; Wheatley Mauche 2005). and the connection between coronae and jets (e.g.. Fender et al.," 1979; Done Osborne 1997; Wheatley Mauche 2005), and the connection between coronae and jets (e.g., Fender et al."
+ 1999: Markolf et al., 1999; Markoff et al.
+ 2005: Rodriguez Prat 2008) has already been proposed., 2005; Rodriguez Prat 2008) has already been proposed.
+ ILowever. (he results of Ishida et al. (," However, the results of Ishida et al. ("
+2009: also LD2008) put the presence of coronae in accretion disks on a solid ground. and further suggest that magnetic activity similar to that in active stars occur in these coronae.,"2009; also LB2008) put the presence of coronae in accretion disks on a solid ground, and further suggest that magnetic activity similar to that in active stars occur in these coronae."
+ The speculative interpretation in (he present letter (Section 2)) is based on three properties of the November 2009 outburst of Aq] X-1 (M2010)., The speculative interpretation in the present letter (Section \ref{sec:flare}) ) is based on three properties of the November 2009 outburst of Aql X-1 (M2010).
+ I emphasize that I do not propose an alternative explanation to (the hard (715 keV) N-rav peak. In N-ray. transients the hard. X-ray. peak can generally be accounted [or bv a disk instability as studied by Dubus et al. (," I emphasize that I do not propose an alternative explanation to the hard $>15 \keV$ ) X-ray peak, In X-ray transients the hard X-ray peak can generally be accounted for by a disk instability as studied by Dubus et al. ("
+2001).,2001).
+ In the present case. the hard X-ray. peak has a triangular shape (see definition in Chen et al.," In the present case, the hard X-ray peak has a triangular shape (see definition in Chen et al."
+ 1997) for ~12 d. The radio peak appears during the decay. phase of the hard. X-ray emission., 1997) for $\sim 12~$ d. The radio peak appears during the decay phase of the hard X-ray emission.
+ The soft (RATE ASAI. 2—10 keV) X-ray. emission appears with the hard. X-ray emission. but ils large and rapid rise starts only alter the radio peak (see below).," The soft (RXTE ASM; $2-10 \keV$ ) X-ray emission appears with the hard X-ray emission, but its large and rapid rise starts only after the radio peak (see below)."
+ The (wo peaks. one in the hard X-rav. followed by one in the soft. X-ray. can be seen for another flare of Aq X-1 in Yu et al. (," The two peaks, one in the hard X-ray followed by one in the soft X-ray, can be seen for another flare of Aql X-1 in Yu et al. ("
+2003). but they do not have radio observations.,"2003), but they do not have radio observations."
+ The two X-ray. peaks do not resemble at all the secondary peaks discussed for N-oray novae by. Chen et al. (, The two X-ray peaks do not resemble at all the secondary peaks discussed for X-ray novae by Chen et al. (
+1997). anc mist be explained bv a dilferent process.,"1997), and must be explained by a different process."
+ Among XRB svstems the double peak structure of the outburst is common. like NTE J18594-226 (Brocksopp et al.," Among XRB systems the double peak structure of the outburst is common, like XTE J1859+226 (Brocksopp et al."
+ 2002) in which the radio peak occurs ab (he start of the extended soft X-ray peak: the hard X-ray peak occurs before the peaks of the radio ancl of the soft X-ray. enissions (Brocksopp οἱ al., 2002) in which the radio peak occurs at the start of the extended soft X-ray peak; the hard X-ray peak occurs before the peaks of the radio and of the soft X-ray emissions (Brocksopp et al.
+ 2002)., 2002).
+ Such a structure is seen in solar flares as well., Such a structure is seen in solar flares as well.
+ In the magnetic flare model the hard X-ray peak is related {ο the event that rapidly amplifies the magnetic field., In the magnetic flare model the hard X-ray peak is related to the event that rapidly amplifies the magnetic field.
+ This field later powers the radio and soft X-raw enilsslons., This field later powers the radio and soft X-ray emissions.
+ The appearance of the hard X-ray. aud soft. X-ray peaks one alter (he other is quite similar (o that seen in DII NRB svstems. but it is not at all similar (o the delay. in rise to maximum between the optical and extveme UV and X-ray emissions in dwarl novae (Mauche el al.," The appearance of the hard X-ray and soft X-ray peaks one after the other is quite similar to that seen in BH XRB systems, but it is not at all similar to the delay in rise to maximum between the optical and extreme UV and X-ray emissions in dwarf novae (Mauche et al."
+ 2001; Wheatley et al., 2001; Wheatley et al.
+ 2003)., 2003).
+ In Aql X-1 the two peaks are separated. while in cdwarl," In Aql X-1 the two peaks are separated, while in dwarf"
+A inajor challenge to the ACDAL cosmological model is its over-prediction of substructure on small scales compared to the actual oserved nunber of dwarf ealaxies in the Local Cacmp (the so-called “uiissing satellite problem: Mooreetal.1999:Ixlvpiuoet 1999)).,"A major challenge to the $\Lambda$ CDM cosmological model is its over-prediction of substructure on small scales compared to the actual observed number of dwarf galaxies in the Local Group (the so-called “missing satellite problem”; \citealt{Moore1999, Klypin1999}) )."
+ The inost welbstudied explanation for this discrepancy is that the lowest mass. halos were inefficient at formuue stars. so that most of these substructures remain dark.," The most well-studied explanation for this discrepancy is that the lowest mass halos were inefficient at forming stars, so that most of these substructures remain dark."
+ Ἱπάσσα. several PPOcesses Call SUpploss star formation aud reduce the umuuber of visible halos. including tidal strippiug «M4 satellites (e.g...Ixavtsovctal. 2001).. supernova feedback. aud the merease of eas temperature due to cosmic reionuization (EfstathiotMadauetal. 2008)..," Indeed, several processes can suppress star formation and reduce the number of visible halos, including tidal stripping of satellites \citep[e.g.,][]{Kravtsov2004b}, supernova feedback, and the increase of gas temperature due to cosmic reionization \citep{Efstathiou1992, Gnedin2000, Benson2002b,
+Somerville2002, Madau2008b}."
+ Dynamical interactions between dwarfs (D'Onghiaetal.20))) or between dwurfs ar the huninous disks of large halos may also play a role., Dynamical interactions between dwarfs \citep{Donghia2009} or between dwarfs and the luminous disks of large halos \citep{Donghia2010} may also play a role.
+" The discovery of a popuation of ""ultra-fünt dwart ealaxies (UFDs: £<LPL. yan the Sloan Dieital Sky Survey (SDSS) (c.e..Willnanetal.2005:Delokurovetal.2006:Zucker2006) has roughlv double the number of known satelites. alleviating the problem somewhat."," The discovery of a population of “ultra-faint” dwarf galaxies (UFDs; $L < 10^5 L_{\odot}$ ) in the Sloan Digital Sky Survey (SDSS) \citep[e.g.,][]{Willman2005, Belokurov2006, Zucker2006b} has roughly doubled the number of known satellites, alleviating the problem somewhat."
+ Based ou the sky coverage aud seusitivitv of SDSS. estimates of a competeness-corrected. Inninositv function down to A42 (Ixoposovetal.2008:Tollerudetal.2008). are avizülable.," Based on the sky coverage and sensitivity of SDSS, estimates of a completeness-corrected luminosity function down to $M_V \sim -2$ \citep{Koposov2008, Tollerud2008} are available."
+ The possible umber of faint saelites could be up to the order of hundreds. but due te» the inconiplee skv coverage of SDSS. the exact nuniCL reniadns uncertain.," The possible number of faint satellites could be up to the order of hundreds, but due to the incomplete sky coverage of SDSS, the exact number remains uncertain."
+ The discovery of these low-huuinositv diwarts las also rise many Low cestions. however.," The discovery of these low-luminosity dwarfs has also raised many new questions, however."
+ Follow-up observatioli have shown that they are the faintest. nost dark luatter-domiunaed and ictal-poor svstenis -iow {51non&Ceha20πιWirbyetal.2008).. with netallicity patterus siniir to the lowest netallicity stars in the halo (c.¢..Fre»etal.BOLO).," Follow-up observations have shown that they are the faintest, most dark matter-dominated and metal-poor systems known \citep{Simon2007, Kirby2008b}, with metallicity patterns similar to the lowest metallicity stars in the halo \citep[e.g.,][]{Frebel2010a}."
+ Iji order to tuuclerstane the cosmological contest of UFDs. including heir relationship to the first galaxies. aud to the building ocks of the stellar halo. it is imiportaut to »» able to uodel conditions in the universe at the time of their orluation (c.@.. Frebel&Broun 2010)).," In order to understand the cosmological context of UFDs, including their relationship to the first galaxies, and to the building blocks of the stellar halo, it is important to be able to model conditions in the universe at the time of their formation (e.g., \citealt{Frebel2010c}) )."
+ Seni-analvtie models of galaxy formation ie been able to explain the observed umber of saelites bv invoking various feedback iuechanisnis. ducudiug an external reionization fold (e.e..Bensonetal.2002):Alaecióetal2010:Lial. 2010).," Semi-analytic models of galaxy formation have been able to explain the observed number of satellites by invoking various feedback mechanisms, including an external reionization field \citep[e.g.,][]{Benson2002a,
+Maccio2010, Li2010}."
+ According to this scenario. as the intergalactie medium ds heated by roeiouization. gas cau no longer be accreted bv aud COO within halos below a certain viria teniperature. inhibiting further star formation iu the lowest mass alos.," According to this scenario, as the intergalactic medium is heated by reionization, gas can no longer be accreted by and cool within halos below a certain virial temperature, inhibiting further star formation in the lowest mass halos."
+ However. these models ecucrally τον ο Qon ΠΠ)lifted. unioru reiouization feld. which is Likely quite dierent from reality.," However, these models generally rely on a simplified, uniform reionization field, which is likely quite different from reality."
+ Simulations of reionization ludicate at it Was likely extended in time aud spatially ithomoge1COlls. wit ra wide ranee im reionization redshifts for dierent pars of the universe," Simulations of reionization indicate that it was likely extended in time and spatially inhomogeneous, with a wide range in reionization redshifts for different parts of the universe"
+range of the observational campaign that discovered it.,range of the observational campaign that discovered it.
+ For objects found in large survevs. the limiting magnitudes are generally well determined.," For objects found in large surveys, the limiting magnitudes are generally well determined."
+ Of the 80 SDOs used here. 31 were discovered by the Deep Ecliptic Survey and 11 were discovered by the Canada-France-IHasvaii Telescope Survey.," Of the 80 SDOs used here, 31 were discovered by the Deep Ecliptic Survey and 11 were discovered by the Canada-France-Hawaii Telescope Survey."
+ The statistics and limiting magnitudes for the Deep Ecliptie Survey and the Canacda-France-Lawai Telescope Survey are described by 2? and? respectively., The statistics and limiting magnitudes for the Deep Ecliptic Survey and the Canada-France-Hawaii Telescope Survey are described by \citet{elliot05} and \citet{trujillo01cfht} respectively.
+ For {he remaining 38 objects not found by dedicated. well characterized Kuiper belt survevs. an effective limiting magnitude was inferred. [rom the distribution of the objects” visual magnitudes upon discovery.," For the remaining 38 objects not found by dedicated, well characterized Kuiper belt surveys, an effective limiting magnitude was inferred from the distribution of the objects' visual magnitudes upon discovery."
+ This limiting magnitude may not be accurate For any individual object. but because the goal of the debiasing procedure is to gain insight to the bulk orbital characteristics of the objects. it is a reasonable approach for this group of objects.," This limiting magnitude may not be accurate for any individual object, but because the goal of the debiasing procedure is to gain insight to the bulk orbital characteristics of the objects, it is a reasonable approach for this group of objects."
+ Of these objects. G were analvzed separately using the same method. because (μον all used data from the Sloan Digital Skv Survey.," Of these objects, 6 were analyzed separately using the same method, because they all used data from the Sloan Digital Sky Survey."
+ Because detecting moving objects was not the primary objective of this survey. characteristics for this use of (he survey have not been published. but analvzing the objects separately vields a better estimate of the limiting magnitude because thev all come from data taken with the same instrument under similar conditions.," Because detecting moving objects was not the primary objective of this survey, characteristics for this use of the survey have not been published, but analyzing the objects separately yields a better estimate of the limiting magnitude because they all come from data taken with the same instrument under similar conditions."
+ To estimate the probability of discovery. each objects orbit was apportioned into Iractions spent within certain ranges of heliocentric distances.," To estimate the probability of discovery, each object's orbit was apportioned into fractions spent within certain ranges of heliocentric distances."
+ The first range of distances corresponds io the range in which the object had effectively a chance of discovery because it would be brighter than the detection elliciency of its discovery group., The first range of distances corresponds to the range in which the object had effectively a chance of discovery because it would be brighter than the detection efficiency of its discovery group.
+ The second range is the range of distances for which the object had less than a chance of discovery. bul more than a chance. corresponding to the discovery. group's limiting magnitude.," The second range is the range of distances for which the object had less than a chance of discovery, but more than a chance, corresponding to the discovery group's limiting magnitude."
+ The fraction of the objects orbit spent outside both these ranges was considered (o have a chance of discovery: we use this simplified scheme because (a) the detection elliciency for any given survey tends to drop off very rapidly with magnitude past the detection limit (see for example ?s Fig., The fraction of the object's orbit spent outside both these ranges was considered to have a chance of discovery; we use this simplified scheme because (a) the detection efficiency for any given survey tends to drop off very rapidly with magnitude past the detection limit (see for example \citet{bernstein04}' 's Fig.
+ 2 or ?s Fig., 2 or \citet{fuentes08}' 's Fig.
+ 1). ancl (b) the uncertainties in the photometry used to determine the visual magnitudes of the discovered objects are high. and thus do not warrant a more elaborate scheme.," 1), and (b) the uncertainties in the photometry used to determine the visual magnitudes of the discovered objects are high, and thus do not warrant a more elaborate scheme."
+ The probability of discovery lor any object is then the fraction of an orbital period the object spends in the first range plus hall the fractional time spent in the second range., The probability of discovery for any object is then the fraction of an orbital period the object spends in the first range plus half the fractional time spent in the second range.
+ This probability can be used to appropriately weight each of the known objects to produce the model population., This probability can be used to appropriately weight each of the known objects to produce the model population.
+ This weighting is included in (ie number of times each of the known SDOs is ‘cloned’ to create our model population., This weighting is included in the number of times each of the known SDOs is `cloned' to create our model population.
+ The nunber of clones is inversely proportional to the detection probability because the detection of objects that spend only a small fraction of their tme in an observable range implies that the intrinsic 5DO population contains many such objects., The number of clones is inversely proportional to the detection probability because the detection of objects that spend only a small fraction of their time in an observable range implies that the intrinsic SDO population contains many such objects.
+ The drawback of this cloning procedure is (hat it leaves some of the observed objects with only one or a [ew clones in a c103 test particle sinuulation., The drawback of this cloning procedure is that it leaves some of the observed objects with only one or a few clones in a $\sim10^3$ test particle simulation.
+ Because (he evolution of any single test particle is not particularly meaningful in a simulation, Because the evolution of any single test particle is not particularly meaningful in a simulation
+Recently. the BLR has received much attention. because the ine widths. in combination with the size of the region. measured or example by reverberation mapping. allows a determination of he mass of the central super-massive black hole (e.g.22).. which is essential if one wants to place AGN in the context of genera galaxy evolution.,"Recently, the BLR has received much attention, because the line widths, in combination with the size of the region, measured for example by reverberation mapping, allows a determination of the mass of the central super-massive black hole \citep[e.g.][]{Benea09,Net09}, which is essential if one wants to place AGN in the context of general galaxy evolution."
+ From such studies. one finds that AGN generally yave normal sized black and that most galaxies must have been active for one or more times in their life.," From such studies, one finds that AGN generally have normal sized black and that most galaxies must have been active for one or more times in their life."
+ The measurement of the black hole mass is affected by the relative importance of radiation pressure anc gravity. Which should dominate the force budget.," The measurement of the black hole mass is affected by the relative importance of radiation pressure and gravity, which should dominate the force budget."
+ This is connected to the properties of individual BLR clouds and the dynamica configuration of the system (2).., This is connected to the properties of individual BLR clouds and the dynamical configuration of the system \citep{Marcea08}.
+ The BLR is a standard AGN ingredient (e.g.222): It is located inside the obscuring torus. above and below the accretion disk. and at a distance of fractions of a parsee from the super-massive black hole.," The BLR is a standard AGN ingredient \citep[e.g.][]{Ost88,Pet97,Net08}: It is located inside the obscuring torus, above and below the accretion disk, and at a distance of fractions of a parsec from the super-massive black hole."
+ The line emission is powered bv photoionisation by a broad-band continuum. due to the innermost parts of the accretion disk., The line emission is powered by photoionisation by a broad-band continuum due to the innermost parts of the accretion disk.
+ Photoionisation also governs the thermodynamics of these clouds: They have a stable equilibrium temperature of order 107 K. independent of their locations.," Photoionisation also governs the thermodynamics of these clouds: They have a stable equilibrium temperature of order $10^4$ K, independent of their locations."
+ The critical densities of observed and suppressed emission lines and pressure equilibrium considerations (2) constrain the number densities in the clouds to n4=10101 u, The critical densities of observed and suppressed emission lines and pressure equilibrium considerations \citep{FE84} constrain the number densities in the clouds to $n_\mathrm{cl}=10^{10\pm1}$ $^{-3}$.
+ Signatures from partially ionised zones are typically observed., Signatures from partially ionised zones are typically observed.
+ Therefore. the Strómmgren depth is a lower limit of the cloud size.," Therefore, the Strömmgren depth is a lower limit of the cloud size."
+ An upper limit may be found by requiring the lines to be smooth. which demands a certain minimum on the number of individual entities.," An upper limit may be found by requiring the lines to be smooth, which demands a certain minimum on the number of individual entities."
+ This results in cloud sizes of about E10727! em therein)., This results in cloud sizes of about $r_\mathrm{cl} = 10^{12\pm1}$ cm .
+ In agreement with this. detailed photoionisation models (2). constrain the column densities to AN.10°? 7.," In agreement with this, detailed photoionisation models \citep{KK81} constrain the column densities to $N_\mathrm{cl}> 10^{22}$ $^{-2}$."
+ Because BLR clouds are optically thick. high ionisation lines are produced only by the illuminated surface of the clouds.," Because BLR clouds are optically thick, high ionisation lines are produced only by the illuminated surface of the clouds."
+ In such lines. one observes generally the far side of the BLR. wherefore outflows manifest themselves as redshifts. inflows as blueshifts.," In such lines, one observes generally the far side of the BLR, wherefore outflows manifest themselves as redshifts, inflows as blueshifts."
+ Tuch information has also been gathered on the dynamical state of the BLR., Much information has also been gathered on the dynamical state of the BLR.
+ Inclination matters: the statistics of line widths excludes Keplerian rotation in a flat disk. but is instead consistent with a thick disk configuration with v/6z2.5. where v is the rotational velocity and o the turbulent one (2)..," Inclination matters: the statistics of line widths excludes Keplerian rotation in a flat disk, but is instead consistent with a thick disk configuration with $v/\sigma \approx 2.5$, where $v$ is the rotational velocity and $\sigma$ the turbulent one \citep{Ost78}."
+ For radio-loud AGN. the ratio of core to lobe power correlates with the width ofthe," For radio-loud AGN, the ratio of core to lobe power correlates with the width ofthe"
+"compute errors on the amplitudes that took into account the correlations between the fitted parameters I,9 and H,o (e.g., ?).","compute errors on the amplitudes that took into account the correlations between the fitted parameters $\Gamma_{n0}$ and $H_{n0}$ \citep[e.g.,][]{Appourcorr}."
+ These amplitudes were then scaled to their bolometric equivalent using the bolometric correction derived from the spectral response of the passband (?)., These amplitudes were then scaled to their bolometric equivalent using the bolometric correction derived from the spectral response of the passband \citep{Ballotbol}.
+" Finally, we obtained a maximum bolometric amplitude of At27470.75ppm (at 833.90 Hz) for KIC 10273246 (cf."," Finally, we obtained a maximum bolometric amplitude of $A_{n0,{\rm{bol}}}^{\rm{(max)}}\!=\!7.17 \pm 0.75\:{\rm{ppm}}$ (at $833.90\:{\rm{\mu Hz}}$ ) for KIC 10273246 (cf."
+" Table[7)), and of AU?)—756+0,99ppm (at 997.14 wHz) for KIC 10920273 (cf."," Table \ref{Freq_Mulder2}) ), and of $A_{n0,{\rm{bol}}}^{\rm{(max)}}\!=\!7.56 \pm 0.99\:{\rm{ppm}}$ (at $997.14\:{\rm{\mu Hz}}$ ) for KIC 10920273 (cf."
+ Table; , Table \ref{Freq_Scully2}) ).
+These values were computed based on the results returnedBp. by the respective and were found to be consistent with the values obtained from the other fitters’ results., These values were computed based on the results returned by the respective and were found to be consistent with the values obtained from the other fitters' results.
+" ? have suggested an empirical scaling relation to predict the amplitudes of solar-like oscillations that, although extensively used, predicts amplitudes in F-type stars that are higher than actually observed (e.g.,?).."," \citet{KB95} have suggested an empirical scaling relation to predict the amplitudes of solar-like oscillations that, although extensively used, predicts amplitudes in F-type stars that are higher than actually observed \citep[e.g.,][]{Michel08}."
+" Recently, the same authors have proposed scaling relation for the amplitudes which is based on simplea newphysical arguments (?):: where L is the stellar luminosity, and the value of r is chosen to be either r=1.5 (assuming adiabatic oscillations) or r—2 (followingafittoobservationaldatain?).."," Recently, the same authors have proposeda new scaling relation for the amplitudes which is based on simple physical arguments \citep{KB11}: : where $L$ is the stellar luminosity, and the value of $r$ is chosen to be either $r\!=\!1.5$ (assuming adiabatic oscillations) or $r\!=\!2$ \citep[following a fit to observational data in][]{KB95}."
+" By assuming that Ymax is a fixed fraction of the acousticcut-off frequency, i.e., and adopting a scaling relation for the stellar mass based on seismic parameters (e.g., ?),, we can combine Eqs. ΠΠ-"," By assuming that $\nu_{\rm{max}}$ is a fixed fraction of the acousticcut-off frequency, i.e., and adopting a scaling relation for the stellar mass based on seismic parameters \cite[e.g.,][]{Kallinger}, , we can combine Eqs. \ref{Abol}-"
+"Π3] to obtain the following relation (normalized with respect to values in the Sun): where An0,bol(mx©=2.53ppm, vao=3050uHz, and Avo=135uHz."," \ref{Mass} to obtain the following relation (normalized with respect to values in the Sun): where $A_{n0,\rm{bol}\,\sun}^{\rm{(max)}}\!=\!2.53\:{\rm{ppm}}$, $\nu_{\rm{max}\,\sun}\!=\!3050\:{\rm{\mu Hz}}$, and $\Delta\nu_{\sun}\!=\!135\:{\rm{\mu Hz}}$."
+" Note that by setting r=1.5 (as we will be assuming hereafter), the dependence of Eq."," Note that by setting $r\!=\!1.5$ (as we will be assuming hereafter), the dependence of Eq."
+ [I4] on Τεῃ is canceled and this scaling relation then solely depends on seismic parameters., \ref{Ascaling} on $T_{\rm{eff}}$ is canceled and this scaling relation then solely depends on seismic parameters.
+" Finally, use of this relation gives a predicted maximum bolometric amplitude of Ao=7.27+1.33ppm for KIC 10273246, and of Aue)=8.24+1.44ppm for KIC 10920273, with the relatively large uncertainties dominated by the errors on Vmax."," Finally, use of this relation gives a predicted maximum bolometric amplitude of $A_{n0,{\rm{bol}}}^{\rm{(max)}}\!=\!7.27 \pm 1.33\:{\rm{ppm}}$ for KIC 10273246, and of $A_{n0,{\rm{bol}}}^{\rm{(max)}}\!=\!8.24 \pm 1.44\:{\rm{ppm}}$ for KIC 10920273, with the relatively large uncertainties dominated by the errors on $\nu_{\rm{max}}$."
+ These values agree with the observed values at the 1-o level., These values agree with the observed values at the $\sigma$ level.
+" This is a particularly interesting result, since no discrepancy is seen between the predicted and observed amplitudes for the F-type star KIC 10273246."," This is a particularly interesting result, since no discrepancy is seen between the predicted and observed amplitudes for the F-type star KIC 10273246."
+" An independent methodology has been used to estimate the large frequency separation and the mode frequencies, which is basedon the analysis of thespectrum with the envelope autocorrelation function (EACF; ?).."," An independent methodology has been used to estimate the large frequency separation and the mode frequencies, which is basedon the analysis of thespectrum with the envelope autocorrelation function \citep[EACF;][]{MA09}. ."
+" As initially proposed by 3, the autocorrelation of the time series — or, equivalently, the"," As initially proposed by \citet{RV06}, , the autocorrelation of the time series – or, equivalently, the"
+to move to the wing sides.,to move to the wing sides.
+" In addition, the profile at z=10.2 shifts to shorter wavelength, and it shows the characteristic shape of gas inflow &Miralda-Escudé"," In addition, the profile at $z=10.2$ shifts to shorter wavelength, and it shows the characteristic shape of gas inflow \citep{Zheng02}."
+" Although our simulation(Zheng includes feedback of stellar 2002)..wind similar to that of Springeletal.(2005b), the Lya line profile indicates gas inflow in the galaxy."," Although our simulation includes feedback of stellar wind similar to that of \citet{Springel05d}, the $\lya$ line profile indicates gas inflow in the galaxy."
+ Our result suggests that high-redshift star-forming galaxies may be fueled by efficient inflow of cold gas from the filaments., Our result suggests that high-redshift star-forming galaxies may be fueled by efficient inflow of cold gas from the filaments.
+ We will study this phenomenon in detail in Yajima et al. (, We will study this phenomenon in detail in Yajima et al. (
+in preparation).,in preparation).
+" On the other hand, it was suggested that the asymmetrically shifted profile to the red wing in some LAEs can be made by outflowing gas distribution (e.g.,Mas-Hesseetal 2003)."," On the other hand, it was suggested that the asymmetrically shifted profile to the red wing in some LAEs can be made by outflowing gas distribution \citep[e.g.,][]{Mas-Hesse03}."
+". The growing hot bubble gas around star-forming region from supernovae or radiative feedback can cause outflowing, neutral gas-shells, which result in red-shifted line profiles."," The growing hot bubble gas around star-forming region from supernovae or radiative feedback can cause outflowing, neutral gas-shells, which result in red-shifted line profiles."
+" Recently, Yamadaetal. observed a sample of 91 LAEs at z—3.1, about (2012)half of which show double peaks of strong-blue and weak-red features thought to be caused by gas outflow, while others show a symmetric single peak in which the flux ratio of blue wing to red one is about unity."," Recently, \citet{Yamada12} observed a sample of 91 LAEs at $z = 3.1$, about half of which show double peaks of strong-blue and weak-red features thought to be caused by gas outflow, while others show a symmetric single peak in which the flux ratio of blue wing to red one is about unity."
+" While our model may explain the latter, the missing outflow features in our line sample is probably due to the limitations of our current simulations, e.g., insufficient spatial resolution and simplified treatment of supernovae feedback."," While our model may explain the latter, the missing outflow features in our line sample is probably due to the limitations of our current simulations, e.g., insufficient spatial resolution and simplified treatment of supernovae feedback."
+" In addition, the line profiles of galaxies at high redshift zZ6 may be highly suppressed and changed by scattering in IGM (e.g.,Santos2004;Dijkstraetal.2007;Zhengetal.2010;Laursen 2011), because the Lyo transmission through IGM is very low at the line center and at shorter wavelengths by the Hubble flow (e.g.,Laursenetal.2011)."," In addition, the line profiles of galaxies at high redshift $z \gtrsim 6$ may be highly suppressed and changed by scattering in IGM \citep[e.g.,][]{Santos04, Dijkstra07, Zheng10, Laursen11}, because the $\lya$ transmission through IGM is very low at the line center and at shorter wavelengths by the Hubble flow \citep[e.g.,][]{Laursen11}."
+". Even at lower redshift z~3, the optical depth of IGM can be high depending on the viewing angle and the location of the galaxy"," Even at lower redshift $z \sim 3$, the optical depth of IGM can be high depending on the viewing angle and the location of the galaxy"
+"see e.g., Min et al.","see e.g., Min et al."
+" 2003, Hony et al."," 2003, Hony et al."
+ 2002) was adopted as the grain density distribution., 2002) was adopted as the grain density distribution.
+" The grain size is such as the volume averaged by a sphere has a radius of 0.05 µπι. For the chemical composition, we used amorphous carbon with the optical properties of Preibisch et al. ("," The grain size is such as the volume averaged by a sphere has a radius of 0.05 $\mu$ m. For the chemical composition, we used amorphous carbon with the optical properties of Preibisch et al. ("
+1993) and a density of 1.85 g cm.,1993) and a density of 1.85 g $^{-3}$.
+ For the external radiation source synthetic models of Aringer et al. (, For the external radiation source synthetic models of Aringer et al. (
+2009) were used.,2009) were used.
+" A grid of attached shells and detached shells was constructed by varying: the temperature of the dust at the inner radius, the optical depth τι at 0.55 um, and the outer radius of the shell."," A grid of attached shells and detached shells was constructed by varying: the temperature of the dust at the inner radius, the optical depth $\mathbf{\tau}_{\lambda}$ at 0.55 $\mu$ m, and the outer radius of the shell."
+ A density power index n = 2 was adopted., A density power index $n$ = 2 was adopted.
+ We first fitted the synthetic to the observed SED and photometry collected from the literature (Fig. 4)), We first fitted the synthetic to the observed SED and photometry collected from the literature (Fig. \ref{SED}) )
+ using a least squares minimisation technique to obtain a rough estimate of the luminosity., using a least squares minimisation technique to obtain a rough estimate of the luminosity.
+ Synthetic intensity brightness maps were then derived from the scaled SED at the same units and wavelengths as the PACS data., Synthetic intensity brightness maps were then derived from the scaled SED at the same units and wavelengths as the PACS data.
+ We then derived the synthetic intensity by integrating the flux in concentric apertures around the star and evaluating the difference between the fluxes in two successive apertures., We then derived the synthetic intensity by integrating the flux in concentric apertures around the star and evaluating the difference between the fluxes in two successive apertures.
+ The synthetic profiles were then combined and fitted to the observed intensity profiles and a best-fit model was selected for each star., The synthetic profiles were then combined and fitted to the observed intensity profiles and a best-fit model was selected for each star.
+" In most cases, more than one detached shell was needed to fit the intensity profile (see Table 2))."," In most cases, more than one detached shell was needed to fit the intensity profile (see Table \ref{t:results}) )."
+ Our preliminary modelling result is shown in Fig. 2.., Our preliminary modelling result is shown in Fig. \ref{profiles}.
+ This could be improved by changing the number of detached shells and/or a modification of the density law., This could be improved by changing the number of detached shells and/or a modification of the density law.
+" Additionally, we assumed that the terminal velocity of the gas is equal to the dust expansion velocity."," Additionally, we assumed that the terminal velocity of the gas is equal to the dust expansion velocity."
+" For AQ And, a value of kkm s! was assumed."," For AQ And, a value of km $^{-1}$ was assumed."
+" This assumption may be unrealistic as different shells for a star may expand at different velocities, but since the mass-loss rate scales linearly with the dust expansion velocity, we can still obtain a rough estimate of the mass-loss rate."," This assumption may be unrealistic as different shells for a star may expand at different velocities, but since the mass-loss rate scales linearly with the dust expansion velocity, we can still obtain a rough estimate of the mass-loss rate."
+" To compute the mass-loss rate, we used the relation from Groenewegen et al. ("," To compute the mass-loss rate, we used the relation from Groenewegen et al. ("
+"1998) applied to the case of a thin shell where M is in Myr1, Va is the dust velocity in km s!, Κ. is the stellar radius in solar radii, rg is the inner dust radius in stellar radii, Q; is the absorption coefficient, a is the dust grain radius in cm, V is the dust-to-gas mass ratio (we assumed 0.005), pg is the grain density in g cm~?, and Y; is the shell's thickness in inner shell units.","1998) applied to the case of a thin shell where $\dot{M}$ is in $_{\sun}$ $^{-1}$, $V_{\rm d}$ is the dust velocity in km $^{-1}$, $R_*$ is the stellar radius in solar radii, $r_{\rm d}$ is the inner dust radius in stellar radii, $Q_{\lambda}$ is the absorption coefficient, $a$ is the dust grain radius in cm, $\Psi$ is the dust-to-gas mass ratio (we assumed 0.005), $\rho_{\rm d}$ is the grain density in g $^{-3}$, and $Y_{\rm d}$ is the shell's thickness in inner shell units."
+ The input for the models and the results are found in Table 2.., The input for the models and the results are found in Table \ref{t:results}.
+" The true mass-loss may be lower, because the present wind velocity is probably lower than that assumed."," The true mass-loss may be lower, because the present wind velocity is probably lower than that assumed."
+" And’’s PACS images show a thin detached shell located at 52"" radius.", 's PACS images show a thin detached shell located at $\arcsec$ radius.
+ In earlier surveys for CO radio line emission AQ And was not detected., In earlier surveys for CO radio line emission AQ And was not detected.
+" A new deep re-observation in CO (1-0) with the mm telescope at Onsala Space Observatory resulted in a non detection (Olofsson, priv."," A new deep re-observation in CO $-$ 0) with the m telescope at Onsala Space Observatory resulted in a non detection (Olofsson, priv."
+ comm.)., comm.).
+" This may be caused by photodissociation, keeping in mind the large physical extension of AQ And's shell and its far distance."," This may be caused by photodissociation, keeping in mind the large physical extension of AQ And's shell and its far distance."
+" Using IRAS data, Young et al. (1993a,, 1993b))"," Using IRAS data, Young et al. \cite{Young93a}, \cite{Young93b}) )"
+ measured a source radius of about 3.3’ but this was only found at 60 uum. An ISO- map at 60 um detected extended emission comparable in size to the ring we see in our data when taking into account ISO-PHOT's lower resolution (online Fig. 3))., measured a source radius of about $\arcmin$ but this was only found at 60 $\mu$ m. An ISO-PHOT map at 60 $\mu$ m detected extended emission comparable in size to the ring we see in our data when taking into account ISO-PHOT's lower resolution (online Fig. \ref{f:AQAndISO}) ).
+ There is no obvious extended emission at 3.3' radius from the star neither in our maps nor in the ISO-PHOT maps., There is no obvious extended emission at $\arcmin$ radius from the star neither in our maps nor in the ISO-PHOT maps.
+ It is possible that the IRAS extended emission is caused by background contamination (CIRR2=3)., It is possible that the IRAS extended emission is caused by background contamination (CIRR2=3).
+ Another possibility is that we filter out the very faint extended emission (see Sect. ??))., Another possibility is that we filter out the very faint extended emission (see Sect. \ref{reduction}) ).
+ The inner part of the ring shows various intensity peaks (Fig. 2))., The inner part of the ring shows various intensity peaks (Fig. \ref{profiles}) ).
+ This points towards mass-loss variation over time., This points towards mass-loss variation over time.
+ Seven models of detached shells with dust temperatures ranging from 28 K to 49 K were needed in order to reproduce the observed intensity profile., Seven models of detached shells with dust temperatures ranging from 28 K to 49 K were needed in order to reproduce the observed intensity profile.
+" This corresponds to a mass-loss variation from 1x107!°Moyr7! for the present mass-loss to 152x107"" Myr! for the older mass-loss (some 0000 years ago).", This corresponds to a mass-loss variation from $\times$ $^{-10}$ $_{\sun} {\rm yr}^{-1}$ for the present mass-loss to $\times$ $^{-7}$ $_{\sun} {\rm yr}^{-1}$ for the older mass-loss (some 000 years ago).
+" From the separation between the intensity peaks, we estimate that the mass-loss varies between every 2000 and 5000 yr."," From the separation between the intensity peaks, we estimate that the mass-loss varies between every 2000 and 5000 yr."
+ Cyg’’s intensity profile is similar to that of AAnd in that both indicate significant variation in the mass-loss rate., 's intensity profile is similar to that of And in that both indicate significant variation in the mass-loss rate.
+" We note that the peak of the outer shell located about 33"" from the central star coincides remarkably well with the CO shell observed by Olofsson et al. (2000)),"," We note that the peak of the outer shell located about $\arcsec$ from the central star coincides remarkably well with the CO shell observed by Olofsson et al. \cite{Olofs00}) ),"
+" though the CO shell is much thinner (2.5"")) than the dust shell (see Fig.1)).", though the CO shell is much thinner ) than the dust shell (see \ref{f:obs}) ).
+ Different scenarios have been proposed to explain the origin of the molecular shell., Different scenarios have been proposed to explain the origin of the molecular shell.
+ Olofsson et al. (1988)), Olofsson et al. \cite{Olofs90}) )
+" suggest a thermal pulse origin during which the star experienced a high mass-loss phase lasting only a few hundred years, 7000 years ago."," suggest a thermal pulse origin during which the star experienced a high mass-loss phase lasting only a few hundred years, 7000 years ago."
+ A wind—wind interaction scenario is proposed by Schóiier et al. (2005)), A wind–wind interaction scenario is proposed by Schöiier et al. \cite{Schoe05}) )
+ and Maercker et al. (, and Maercker et al. (
+2010).,2010).
+" In this case, a rapidly moving shell is colliding with a slower moving wind sweeping up old material."," In this case, a rapidly moving shell is colliding with a slower moving wind sweeping up old material."
+ Another possible explanation mentioned by Wareing et al. (, Another possible explanation mentioned by Wareing et al. (
+2006) is a wind-ISM interaction.,2006) is a wind–ISM interaction.
+" For this to happen, the star's space velocity with respect to the ISM has to be high enough to induce a shock with the ISM."," For this to happen, the star's space velocity with respect to the ISM has to be high enough to induce a shock with the ISM."
+ Bow shocks have been observed for AGB stars such as for R Hya (Ueta et al., Bow shocks have been observed for AGB stars such as for R Hya (Ueta et al.
+ 2006) and with PACS data for CW Leo (Ladjal et al., 2006) and with PACS data for CW Leo (Ladjal et al.
+ 2010)., 2010).
+Of the possible nebular types considered the simplest is à conventional LILLE region. photoionised byT317.,"Of the possible nebular types considered the simplest is a conventional HII region, photoionised by."
+. A second. possibility is that the nebula is an ΟΛ., A second possibility is that the nebula is an SNR.
+ bo view of the association with the BeNRB system. it would be extremely interesting to identify the SNR. corresponding to the formation of the svstem's neutron star.," In view of the association with the BeXRB system, it would be extremely interesting to identify the SNR corresponding to the formation of the system's neutron star."
+ This would. amongst other things. enable the system's age to be etermined.," This would, amongst other things, enable the system's age to be determined."
+ Another possibility which exists is that the nebulosity is à bowshock., Another possibility which exists is that the nebulosity is a bowshock.
+ Such structures arise when ran. pressure ue to supersonic motion of a star confines its expanding wind., Such structures arise when ram pressure due to supersonic motion of a star confines its expanding wind.
+ The head or apex of the structure delines a point where the momentum of the expanding stellar wind balances that of the oncoming interstellar medium., The head or apex of the structure defines a point where the momentum of the expanding stellar wind balances that of the oncoming interstellar medium.
+ Bowshocks are commonplace around OB runaway stars., Bowshocks are commonplace around OB runaway stars.
+ Indeed. according to the standard theory of BeXNRB formation. such systems ADJUST be runaways (vandenLleuveletal.2000).," Indeed, according to the standard theory of BeXRB formation, such systems MUST be runaways \cite{vandenheuvel2000}."
+.. Pherefore bowshocks could. be commonplace around ον»., Therefore bowshocks could be commonplace around BeXRBs.
+ The nature of the nebula is discussed below., The nature of the nebula is discussed below.
+ No nebulosity had. been observed when the SAO spectra were taken so the exposures were calculated to. properly expose for the De star., No nebulosity had been observed when the SAAO spectra were taken so the exposures were calculated to properly expose for the Be star.
+ Llowever since the slit of the spectrograph was aligned E-W subsequent. examination of the raw spectral images revealed weak nebular emission lines olfset to one side of the stellar continuum., However since the slit of the spectrograph was aligned E-W subsequent examination of the raw spectral images revealed weak nebular emission lines offset to one side of the stellar continuum.
+ This one-sidedness stems from the asvmametrical morphology of the nebulosity., This one-sidedness stems from the asymmetrical morphology of the nebulosity.
+ Two spectra of the source were acquired. the second because the intended star largely missed. the slit on the first.," Two spectra of the source were acquired, the second because the intended star largely missed the slit on the first."
+ However both contained strong nebular features. ancl so were co-added.," However both contained strong nebular features, and so were co-added."
+ The exposures were sequential and as a pair were bracket(ος by arc exposures: no wavelength shift was detectable., The exposures were sequential and as a pair were bracketted by arc exposures; no wavelength shift was detectable.
+ ALL further reduction took place on this co-acded image., All further reduction took place on this co-added image.
+ The SAAO blue exposure similarly shows only very weak nebular emission., The SAAO blue exposure similarly shows only very weak nebular emission.
+ In view of the existence of a [ar higher S/N blue spectrum from La Silla. no further work was undertaken on the blue SAAO spectrum.," In view of the existence of a far higher S/N blue spectrum from La Silla, no further work was undertaken on the blue SAAO spectrum."
+ Thus the task was used. to extract the region immediately adjacent to the star including these emission features: sky subtraction was from a region 10 pixels below., Thus the task was used to extract the region immediately adjacent to the star including these emission features; sky subtraction was from a region 10 pixels below.
+ Whilst sky subtraction was used to identify and reject. telluric: ancl other. non-nebular lines. actual measurement used. subtraction of a constant. background to significantly reduce noise.," Whilst sky subtraction was used to identify and reject telluric and other non-nebular lines, actual measurement used subtraction of a constant background to significantly reduce noise."
+ This also enabled. telluric lines to be used as wavelength standards. for example the separation of telluric from its SAIC component provides a highly accurate velocity.," This also enabled telluric lines to be used as wavelength standards, for example the separation of telluric from its SMC component provides a highly accurate velocity."
+ Care was taken to exclude anv stellar (including Be circumstellar clise) Lux., Care was taken to exclude any stellar (including Be circumstellar disc) flux.
+ ‘Table 2. lists the nebular emission lines obtained from the SAAO red and ESO blue spectra.," Table \ref{tab:nebulalinelist}
+ lists the nebular emission lines obtained from the SAAO red and ESO blue spectra."
+ The ESO blue spectra taken on 3 November 1999 were extracted using the same method., The ESO blue spectra taken on 3 November 1999 were extracted using the same method.
+ Dividing the lux density integrated along the slit by the approximate scale of for the object one derives a Dux density of ~2010οςοι/sfaresec2 which is close to the value found for the bowshock associated with Vela N-1. (Ixaperctal.1997). of 101erg£cm/sfaresec.zE ," Dividing the flux density integrated along the slit by the approximate scale of for the object one derives a flux density of $\sim20\times10^{-16} erg/cm^2/s/arcsec^2$, which is close to the value found for the bowshock associated with Vela X-1 \cite{kaper97} of $\sim10\times10^{-16} erg/cm^2/s/arcsec^2$ ."
+Decause of the small number of counts from the nebular lines. the extracted spectra are a spatial average over all the nebulosity which fell upon the spectrograph slit.," Because of the small number of counts from the nebular lines, the extracted spectra are a spatial average over all the nebulosity which fell upon the spectrograph slit."
+ Without higherg S/N data this is the only way of extractingSs spectra of usable quality., Without higher S/N data this is the only way of extracting spectra of usable quality.
+ Discussion on the spatial distribution of the different emission lines is presented below., Discussion on the spatial distribution of the different emission lines is presented below.
+ In order to derive the most accurate line [uxes the Emission Line Fitting (ELE) suite within DIPSO was emploved., In order to derive the most accurate line fluxes the Emission Line Fitting (ELF) suite within DIPSO was employed.
+ The JFWNILLM. values were specified. from. measurement of the line to reduce the number of (ree parameters and thus errors., The FWHM values were specified from measurement of the line to reduce the number of free parameters and thus errors.
+ The final fit is shown in Figure 9.., The final fit is shown in Figure \ref{fig:OIIfit}.
+ The velocities are. derived. from. the shift. relative to laboratory rest wavelengths. and further corrected. to a heliocentric reference frame using the program RV.," The velocities are derived from the shift relative to laboratory rest wavelengths, and further corrected to a heliocentric reference frame using the program RV."
+ Phe positive values refer to à movement away from the observer. Le. a recshitt.," The positive values refer to a movement away from the observer, i.e. a redshift."
+ The presence of the stellar continuum alongside the nebular spectrum allows an absolute Iux calibration to be performed., The presence of the stellar continuum alongside the nebular spectrum allows an absolute flux calibration to be performed.
+ Such calibration is important. lor establishing reliable line ratios for diagnostic techniques., Such calibration is important for establishing reliable line ratios for diagnostic techniques.
+ The reddened. theoretical Ixurucz model photospheric spectrum (Ixurucz.1979) fitted. by CHI. was. further reddened to allow for the cireumstellar reddening., The reddened theoretical Kurucz model photospheric spectrum \cite{kurucz1979} fitted by CHR was further reddened to allow for the circumstellar reddening.
+ This was normalized using the photometry at Band Ro for the blue ancl red spectra respectively., This was normalized using the photometry at B and R for the blue and red spectra respectively.
+ The extracted stellar spectrum was divided into this and fitted with a polvnomial vielding the required. sealing factor at cach wavelength to convert [rom counts (the units of the line fitting) to ergf/em?fsfol., The extracted stellar spectrum was divided into this and fitted with a polynomial yielding the required scaling factor at each wavelength to convert from counts (the units of the line fitting) to $erg/cm^2/s/\AA$.
+ Clearly this calibration technique is dependent upon the ]|xurucz atmosphere. being an accurate model of the true stellar spectrum., Clearly this calibration technique is dependent upon the Kurucz atmosphere being an accurate model of the true stellar spectrum.
+ Note however that only the shape of, Note however that only the shape of
+In principle. one can consider that because of a certain mechanism (e.g.. dissipation of mechanical energy). the outer parts of the present V838 Mon wind become excited. producing the observed II] emission lines.,"In principle, one can consider that because of a certain mechanism (e.g., dissipation of mechanical energy), the outer parts of the present V838 Mon wind become excited, producing the observed ] emission lines."
+ In this case. however. the radial velocity of the lines would be close to that of V838 Mon.," In this case, however, the radial velocity of the lines would be close to that of V838 Mon."
+ This is certainlly not the case., This is certainlly not the case.
+ The II] lines have a radial velocity of 13.3+0.7 kkmss! (Sect., The ] lines have a radial velocity of $13.3 \pm 0.7$ $^{-1}$ (Sect.
+ 4.1.1 in Paper D. while the radial velocity of V838 Mon is at least 53 kkmss7! (see Sect. ?22)).," 4.1.1 in Paper I), while the radial velocity of V838 Mon is at least $53$ $^{-1}$ (see Sect. \ref{v_rad}) )."
+ In this case. it would also be difficult to explain the observed II] profiles.," In this case, it would also be difficult to explain the observed ] profiles."
+ All the observed characteristics of the IL] emission lines can. however. be easily explained. 1f they are assumed to be produced by the matter ejected during the 2002 eruption that approaches the B3V companion.," All the observed characteristics of the ] emission lines can, however, be easily explained, if they are assumed to be produced by the matter ejected during the 2002 eruption that approaches the B3V companion."
+ As shown in Sect. 4..," As shown in Sect. \ref{feII_profil},"
+ the observed line profiles as well as their radial velocities can then be well accounted for., the observed line profiles as well as their radial velocities can then be well accounted for.
+ The continuous strengthening of the lines in 2005-2000 is also easy to understand: larger and larger amounts of matter become excited as it approaches the source of excitation., The continuous strengthening of the lines in 2005–2006 is also easy to understand: larger and larger amounts of matter become excited as it approaches the source of excitation.
+ Excitation by radiation from the B3V companion then explains why the emission line spectrum is dominated byFell., Excitation by radiation from the B3V companion then explains why the emission line spectrum is dominated by.
+ In the spectrum of a B3V star. there are enough photons capable of ronizing species with an ionization potential of ~7 eV. e.g.. Fe. Ni. while there are very few photons capable of ronizing the most abundant elements such as H. He or CNO.," In the spectrum of a B3V star, there are enough photons capable of ionizing species with an ionization potential of $\sim 7$ eV, e.g., Fe, Ni, while there are very few photons capable of ionizing the most abundant elements such as H, He or CNO."
+ Finally. the idea easy explains the ~70 day eclipse-like event observed in November/December 2006 as an occultation by a dense cloud of the matter ejected in 2002 and now crossing the line of sight of the B3V companion.," Finally, the idea easy explains the $\sim 70$ day eclipse-like event observed in November/December 2006 as an occultation by a dense cloud of the matter ejected in 2002 and now crossing the line of sight of the B3V companion."
+ If this is the case. one can expect that the 2006 eclipse should be followed by similar events when other fragments of the 2002 ejecta cross the line of sight.," If this is the case, one can expect that the 2006 eclipse should be followed by similar events when other fragments of the 2002 ejecta cross the line of sight."
+ Indeed. a few months later V838 Mon entered another eclipse-like We note that the latter would be. at least. difficult to explain as a phenomenon typical of an eclipsing binary system.," Indeed, a few months later V838 Mon entered another eclipse-like We note that the latter would be, at least, difficult to explain as a phenomenon typical of an eclipsing binary system."
+ Several molecules observed in the spectrum of V838 Mon reveal bands formed from high excitation-energy levels (see Sect., Several molecules observed in the spectrum of V838 Mon reveal bands formed from high excitation-energy levels (see Sect.
+ 5 in Paper D)., 5 in Paper I).
+ They are usually indicative of a radial heliocentric velocity of 58+5 ((see also Sect. 5.1))., They are usually indicative of a radial heliocentric velocity of $58\pm5$ (see also Sect. \ref{bands}) ).
+ It is reasonable to assume that they arise very close to the photosphere and have a velocity close to the stellar one., It is reasonable to assume that they arise very close to the photosphere and have a velocity close to the stellar one.
+ Earlier determinations of the stellar velocity in Kolevetal.(2002) and Kipperetal.(2004) implied a value ofkms.. Le.. close to our result.," Earlier determinations of the stellar velocity in \cite{kolev} and \cite{kipp04} implied a value of, i.e., close to our result."
+ However. the SiO maser emission observed in V838 Mon (Deguchietal..2005;Claussen2009) is at a heliocentric velocity ofkms.," However, the SiO maser emission observed in V838 Mon \citep{degu,clauss,degkam} is at a heliocentric velocity of."
+. In the case of late-type stars. the SiO masers usually have a reliable measure of their stellar radial velocity.," In the case of late-type stars, the SiO masers usually have a reliable measure of their stellar radial velocity."
+ There are two possible explanations of this diserepaney., There are two possible explanations of this discrepancy.
+ If the SiO maser emission really originates at the stellar velocity then the estimate made in Paper | may indicate that we do not measure molecular bands deep enough that are to reach hydrostatic levels of the V838 Mon atmosphere. re.. even the highest excitation bands are formed in already expanding layers.," If the SiO maser emission really originates at the stellar velocity then the estimate made in Paper I may indicate that we do not measure molecular bands deep enough that are to reach hydrostatic levels of the V838 Mon atmosphere, i.e., even the highest excitation bands are formed in already expanding layers."
+ The other possibility 1s that the optical spectroscopic studies provide à correct estimate of the stellar velocity but that the SiO maser is not at the stellar velocity., The other possibility is that the optical spectroscopic studies provide a correct estimate of the stellar velocity but that the SiO maser is not at the stellar velocity.
+ We note that the radial velocity and width of the main component of the SiO maser are very close to those of the CO rotational lines observed in diffuse matter in the close vicinity of V838 Mon (Kaminski.2008)., We note that the radial velocity and width of the main component of the SiO maser are very close to those of the CO rotational lines observed in diffuse matter in the close vicinity of V838 Mon \citep{kam08}.
+. This point certainly deserves further investigation., This point certainly deserves further investigation.
+ In the present paper. we usually adopt the radial velocity derived from the SiO maser. i.e..kms.. as the radial velocity of V838 Mon.," In the present paper, we usually adopt the radial velocity derived from the SiO maser, i.e., as the radial velocity of V838 Mon."
+ One has. however. to keep in mind that the results from optical spectroscopy infer a value lower.," One has, however, to keep in mind that the results from optical spectroscopy infer a value lower."
+ As discussed in Sect. 2..," As discussed in Sect. \ref{decline_sect},"
+ we assume that the [FeIL] emission lines are formed in the matter ejected by V838 Mon during its 2002 outburst. which is now approaching the B3V companion and ionized by the UV radiation of the companion.," we assume that the ] emission lines are formed in the matter ejected by V838 Mon during its 2002 outburst, which is now approaching the B3V companion and ionized by the UV radiation of the companion."
+" In this section, using simplified modelling. we show that the observed profiles of the [Fell] lines can be easily explained by this scenario."," In this section, using simplified modelling, we show that the observed profiles of the ] lines can be easily explained by this scenario."
+ We consider a system of two stars. S$) and S». with a separation A (see Fig. 1)).," 	 We consider a system of two stars, $S_1$ and $S_2$, with a separation $A$ (see Fig. \ref{sketch}) )."
+ We assume that the star ο is a source of a steady. spherically symmetric wind. that forms," We assume that the star $S_1$ is a source of a steady, spherically symmetric wind, that forms"
+during November 5-7 is ~—L2 but on December 13 PY appears to go below1.,during November 5-7 is $\sim -1.2\%$ but on December 13 $P\%$ appears to go below.
+54... These values are close to the values generally. observed in other comets at. such low phase angles. eg.," These values are close to the values generally observed in other comets at such low phase angles, eg."
+ comet LHallev. (Senetal.1991).. Lale-Bopp (Ganeshetal.1998)... LINEAR WALL (Joshietal.2002. 2003).," comet Halley \citep{Sen1991}, Hale-Bopp \citep {Ganesh1998}, LINEAR WM1 \citep{Joshi2002,Joshi2003}."
+. The data points Lic close to the best Ge vs a curve). to further supporting the argument that το/LHolmes is a usual comet.," The data points lie close to the best fit vs $\alpha$ curve), to further supporting the argument that 17P/Holmes is a usual comet."
+ On the other hand. Rosenbushetal.(2009)— have made almost all their observations in. broad bands(WRC. Lh. D.," On the other hand, \cite{rosenbush2009} have made almost all their observations in broad bands(WRC, R, I)."
+" Their only observation in BC ""narrow! band on October 27. 2007 gives polarisation. value which is very close. (within 10) to the typical pmse-polarisation curve for comets(cf Figure 2)) and also agree with polarisation values. obtained. by us ab later ¢late."," Their only observation in BC 'narrow' band on October 27, 2007 gives polarisation value which is very close (within $1\sigma$ ) to the typical phase-polarisation curve for comets(cf Figure \ref{PP6840}) ) and also agree with polarisation values obtained by us at later date."
+ Lt is. therefore. very likely that the broad band coninuum Lux is contaminated. by the eas emission. lowering he absolute value of the degree of polarisation.," It is, therefore, very likely that the broad band continuum flux is contaminated by the gas emission, lowering the absolute value of the degree of polarisation."
+ Strong gas emission has been detected during the outburst of 17P/Holmes (e.g.Bockelée-Morvan:20082.b:Averοἱal.2008:Scheicher2007 ).. which supports the argument given above for the lower absolute value. of. the polarisation," Strong gas emission has been detected during the outburst of 17P/Holmes \citep[e.g.][]{bockelee2008a,
+bockelee2008apj, biver2008, schleicher2007}, which supports the argument given above for the lower absolute value of the polarisation"
+παται latitudes for photons >100 MeV: DMM2.0«107slILatom (Abdoetal.2009b).,"medium latitudes for photons $> 100$ MeV: $\Gamma_{\pi \rightarrow \gamma\gamma}^{\rm MW} = 2.0 \times 10^{-25} \
+ \rm s^{-1} \mbox{H-atom}^{-1}$ \citep{FermiMW09}."
+. This value represeuts a large-scale (~kpc) spatial averacing of Galactic cosmic-ray properties. which is appropriate for our elobal calculation.," This value represents a large-scale $\sim$ kpc) spatial averaging of Galactic cosmic-ray properties, which is appropriate for our global calculation."
+" We can now write the exnmnierav hmuuinositv deusitv for normal galaxies j= esfr((.}(13) in terms of the cosmic star-formation rate p,(:) fOr."," We can now write the gamma-ray luminosity density for normal galaxies ,z) = (z) in terms of the cosmic star-formation rate $\csfr(z) = \avg{\psi n_{\rm gal}}$ ."
+" We define a mean interstellar gas 1iass as The assumption that losses are ""ENescape-domiuated(10 aud nuiform across galaxies cau only be approximately valid at best.", We define a mean interstellar gas mass as = The assumption that losses are escape-dominated and uniform across galaxies can only be approximately valid at best.
+ This is a major uncertaiuty in our model. which will benefit from future data ou the EGB aud resolved ealaxies.," This is a major uncertainty in our model, which will benefit from future data on the EGB and resolved galaxies."
+ For example. eveu the leakw-box model can eeneralize eq. (2))," For example, even the leaky-box model can generalize eq. \ref{eq:scaling}) )"
+ to L.xΑιdea: Variations in the enerey dependence of the escape leneth X would also change the cosiuc-orav and photon spectral indices which we take as universal., to $L_\gamma \propto \Lambda_{\rm esc} \psi \mgas$; variations in the energy dependence of the escape length $\Lambda_{\rm esc}$ would also change the cosmic-ray and photon spectral indices which we take as universal.
+ Tudeed.Ferm: observations of the Large.Masellauic Cloud suggest that cosmnüc-cav coufinement aud propagation are non-trivial (Abdo—ctal2010€).," Indeed, observations of the LargeMagellanic Cloud suggest that cosmic-ray confinement and propagation are non-trivial \citep{fermi-LMC}."
+ Moreover. starburst galaxies show very high cosmic-ray iuteusities within simall volumes where inclastic collisions conrpete with. and sometimes donunate. outflows to regulate cosmic-ray losses (Paghoueetal.1996:Lac 2007).," Moreover, starburst galaxies show very high cosmic-ray intensities within small volumes where inelastic collisions compete with, and sometimes dominate, outflows to regulate cosmic-ray losses \citep{pag1996,lacki,torres,thompson,persic,stecker07}."
+. For this reason. eq. (9))," For this reason, eq. \ref{eq:gammalum}) )"
+ provides a rough description of normal escape-doninated galaxies only: we do not expect it to hold for starburst galaxies. which we will exclude below.," provides a rough description of normal escape-dominated galaxies only; we do not expect it to hold for starburst galaxies, which we will exclude below."
+ We can infer a galaxv interstellar eas mass af a eiven star-formation rate.s via the well-established Izenuicutt-Schiidt law (Schuuidt1959:Kennicutt1998).," We can infer a galaxy's interstellar gas mass at a given star-formation rate, via the well-established Kennicutt-Schmidt law \citep{schmidt,kennicutt}."
+". The for star formation and gas are fouud to be correlated via X,/M.xr!kpe?=(2.540.7)«10HXM.pe2)"" with c=LL015."," The for star formation and gas are found to be correlated via $\dot{\Sigma}_{\star}/\msol \rm yr^{-1} kpc^{-2} = 
+(2.5 \pm 0.7) \times 10^{-4} 
+ (\Sigma_{\rm gas}/\msol {\rm pc}^{-2})^{\mathnormal x}$ with $x = 1.4 \pm 0.15$."
+" Dot[um normal and starburst galaxies follow this correlation. butnormal galaxies populate X,ioacoPaxSenorial_0.1M.vrtkpeD7 whileB starbursts occupy the opposite- reeime (Ikennicutt1998)."," Both normal and starburst galaxies follow this correlation, butnormal galaxies populate $\dot{\Sigma}_{\star,\rm normal} \la \dot{\Sigma}_{\star,\rm normal}^{\rm max} \equiv 0.4 \ \msol\ \rm yr^{-1} kpc^{-2}$ while starbursts occupy the opposite regime \citep{kennicutt}."
+. To recover a relationship between the elobal ealactic star-formation rate ο=Ts and gas mass Aa=RU requires a galactie hiedisks scale leneth raids.," To recover a relationship between the global galactic star-formation rate $\psi = \pi r_{\rm disk}^2 \dot{\Sigma}_{\star}$ and gas mass $\mgas = \pi r_{\rm disk}^2 \Sigma_{\rm gas}$ , requires a galactic disks scale length $r_{\rm disk}$."
+ We take Mens:rag=189kpe/(l|2): observations (Erbetal.2006) indicate that this choice is uncorrelated with ealaxw mass and star-formation rate at a fixed 2.," We take $ r_{\rm disk} = 18.9 \ {\rm kpc}/(1+z)$; observations \citep{erb06}
+ indicate that this choice is uncorrelated with galaxy mass and star-formation rate at a fixed $z$."
+ Combining this with the Nenuicutt-Schinidt law. we fiud τοις1*(15) where ο=δαLor)0.571 and w=έςTLL.," Combining this with the Kennicutt-Schmidt law, we find ,z) = 2.8 where $\beta = 2(1-1/x) = 0.571$ and $\omega = 1/x = 0.714$."
+ Iu our nodel. normal galaxies extend to :) πι.2 ντ this cutoff becomes iuportaut for +z1.Combining eqs. (2.1))," In our model, normal galaxies extend to (z) (z)^2 yr this cutoff becomes important for $z \ga 1$.Combining eqs. \ref{eq:SK-global}) )"
+ aud (9)). we can express a ealaxyvs eanuna-rav bhuuinositv m terms of its SFR aud redshift: Loo.2) Available data ou resolved i;2:0 star-forming ealaxies are consistent with this scaling 2010a)..," and \ref{eq:gammalum}) ), we can express a galaxy's gamma-ray luminosity in terms of its SFR and redshift: ,z) Available data on resolved $z\approx 0$ star-forming galaxies are consistent with this scaling \citep{sfgal}."
+" A ealaxvs huninosity in the ]line provides a well-established tracer of star formation rate: Ly,/1.26«41071W—eftΑνν| (Hoplius2001)..."," A galaxy's luminosity in the line provides a well-established tracer of star formation rate: $\lha/1.26 \times 10^{34} \ {\rm W} 
+= \psi/1 \ \msol \, {\rm yr^{-1}}$ \citep{hopkins}."
+" The dadistribution of galaxies (Iuniuosity function) is related to the cosmic star-formation rate densitv via p,(:)(onxfEg.n(Egu.:) dL, which by now is as a function of redshift both by aand by other methods. so that its shape is well-eternuned observationally."," The distribution of galaxies (luminosity function) is related to the cosmic star-formation rate density via $\csfr(z) = \avg{\psi n} \propto \int \lha \ n(\lha,z) \ d\lha$ , which by now is well-measured as a function of redshift both by and by other methods, so that its shape is well-determined observationally."
+" The eanuua-ray dundünositv density is a differcut noment of the iuumnositv function£C.(2)=(nmX(1l|2)squeLuoμμ.)dL, via our scaling aws above."," The gamma-ray luminosity density is a different moment of the luminosity function$\grld(z) = \avg{\psi n} \propto (1+z)^{-\beta} \int^{\lha^{\rm max}} \lha^{\omega} \ n(\lha,z) \ d\lha$ via our scaling laws above."
+ We see that the eanumuia-rayv απο]τν istribution (ie. lLunuimositv function) at a eiven redshift follows directly from the distribution of stm-Oratfion rates. as traced by Π," We see that the gamma-ray luminosity distribution (i.e., luminosity function) at a given redshift follows directly from the distribution of star-formation rates, as traced by ."
+αν Since we consider rere ouly normal galaxies. we include ouly ealaxies with COXore (00. 2.1) ," Since we consider here only normal galaxies, we include only galaxies with $\psi \le \psi_{\rm normal}^{\rm max}$ (eq. \ref{eq:max-sfr}) )"
+"which sets correspouding lanits Ly,€LYS. aud L.€LUUS* in eq. (2))."," which sets corresponding limits $\lha \le \lha^{\rm max}$, and $L_\gamma \le L_\gamma^{\rm max}$ in eq. \ref{eq:grld}) )."
+" Cinrent data ou the Ihuninositv function can be fit to à Schechter function. of the fon o(L.:)dL=n.(LfEL.)""exptL/L.)L where Lo—πω."," Current data on the luminosity function can be fit to a Schechter function, of the form $n(\tracer,z) \ d\tracer 
+ = n_* \ (\tracer/\tracer_*)^{-\alpha} \
+ \exp\left( - \tracer/\tracer_* \right) \
+ d\tracer/\tracer_*$ where $\tracer = \lha$."
+ Present data are consistent with the value a=1.13. found for :=0. persisting for all redshitts.," Present data are consistent with the value $\alpha = 1.43$, found for $z=0$, persisting for all redshifts."
+ Data also fix the +=0 values Ην)=1.0s10°Mpe aud £.(0)=9.5«10W. for ho=0.71 (Nakamura®ctal.2001).," Data also fix the $z=0$ values $n_*(0) = 1.0 \times 10^{-3} \ {\rm Mpc^{-3}}$ and $\tracer_*(0) = 9.5 \times 10^{34} \ \rm W$, for $h = 0.71$ \citep{nakamura}."
+ This corresponds to a star- rate c(L.)=7.5AL.ντ., This corresponds to a star-formation rate $\psi(\tracer_*) = 7.5 \ \msol/{\rm yr}$.
+" However. observations currently do nof give ""nanbieuous solutions for the other two parameters.— the characteristic comoving density of star-forming galaxies Ds aud the characteristic Ihinunesitv L. (Hopkius 2001)."," However, observations currently do not give unambiguous solutions for the other two parameters, the characteristic comoving density of star-forming galaxies $n_*$ and the characteristic luminosity $\tracer_*$ \citep{hopkins}. ."
+. Moreover. it is unclear whether and how cach parameter with redshift.," Moreover, it is unclear whether and how each parameter with redshift."
+ Two liniting casesbracket the possible behaviors of the Ihuninositv function aud thus of cosnic star formation., Two limiting casesbracket the possible behaviors of the luminosity function and thus of cosmic star formation.
+ Iu the case of the comoving density of stars is fixed. ην=const independent of :. ancl all redshift evolution lies in £.= £..(2).," In the case of the comoving density of stars is fixed, $n_* = const$ independent of $z$ , and all redshift evolution lies in $\tracer_* = \tracer_*(z)$ ."
+ Couverscly. places the redshift evolution iu the density scale 042) while setting £.= const. In cach of," Conversely, places the redshift evolution in the density scale $n_*(z)$ while setting $\tracer_* = const$ In each of"
+The last twenty years have witnessed a significant improvement in the caleulation. of the relevant input physics for stellar evolution calculations. like the equation of state of the stellar matter (Le. Rogers. Swenson Iglesias 1996). radiative opacities (1.e.. Iglesias Rogers 1996. Ferguson et al.,"The last twenty years have witnessed a significant improvement in the calculation of the relevant input physics for stellar evolution calculations, like the equation of state of the stellar matter (i.e., Rogers, Swenson Iglesias 1996), radiative opacities (i.e., Iglesias Rogers 1996, Ferguson et al."
+ 2005). nuclear cross sections (ie.. Angulo et al.," 2005), nuclear cross sections (i.e., Angulo et al."
+ 1999). neutrino emission rates (1.e.. Haft. Raffelt. Weiss 1994).," 1999), neutrino emission rates (i.e., Haft, Raffelt, Weiss 1994)."
+ Several of these advances have been stimulated by the need to match the high-precision data from helioseismology. and an immediate consequence has been the improved accuracy of low mass-main sequence (MS) models for both Population | and II stars.," Several of these advances have been stimulated by the need to match the high-precision data from helioseismology, and an immediate consequence has been the improved accuracy of low mass-main sequence (MS) models for both Population I and II stars."
+ A set of key ingredients in the calculation of stellar models are the nuclear reaction rates. and a large effort has been devoted to improve the measurements of rates at energies as close as possible to the Gamow peak. Le. the energies at which nuclear reactions occur in stars.," A set of key ingredients in the calculation of stellar models are the nuclear reaction rates, and a large effort has been devoted to improve the measurements of rates at energies as close as possible to the Gamow peak, i.e. the energies at which nuclear reactions occur in stars."
+ Thanks to these studies the reaction rates involved in the chain have nowadays a small uncertainty. and the uncertainty on the predicted age - luminosity relationship for low-mass MS stars is also negligible (<2% - Chaboyer et al.," Thanks to these studies the reaction rates involved in the chain have nowadays a small uncertainty, and the uncertainty on the predicted age - luminosity relationship for low-mass MS stars is also negligible $<$ - Chaboyer et al."
+ 1998. Brocato et al.," 1998, Brocato et al."
+ 1998)., 1998).
+ However. near the end of the MS -- e.g.. when the abundance of protons in the core is approaching zero — the energy supplied by the chain becomes insufficient and the star reacts by contracting its core to release gravitational energy.," However, near the end of the MS – e.g., when the abundance of protons in the core is approaching zero – the energy supplied by the chain becomes insufficient and the star reacts by contracting its core to release gravitational energy."
+ As à consequence. according to the virial theorem. both central temperature and density increase and. when the temperature attains a value of ~15x10°K. the H-burning process becomes controlled by the cycle. whose efficiency is critically dependent on the UN(p.y)O reaction rate. the slowest reaction of the whole cycle.," As a consequence, according to the virial theorem, both central temperature and density increase and, when the temperature attains a value of $\sim15\times10^6K$, the H-burning process becomes controlled by the cycle, whose efficiency is critically dependent on the $^{14}N(p,\gamma)^{15}O$ reaction rate, the slowest reaction of the whole cycle."
+" Until few years ago. the rate for the ""N(p.y)?O reaction was uncertain by a factor of 5 at least. because all available laboratory measurements were performed at energies well above the range of interest for astrophysical purposes (Angulo et al."," Until few years ago, the rate for the $^{14}N(p,\gamma)^{15}O$ reaction was uncertain by a factor of 5 at least, because all available laboratory measurements were performed at energies well above the range of interest for astrophysical purposes (Angulo et al."
+ 1999)., 1999).
+ Recently. the LUNA experiment (Formicola et al.," Recently, the LUNA experiment (Formicola et al."
+ 2004) has significantly improved the low energy measurements. obtaining an estimate which is about a factor of 2 lower than previous determinations.," 2004) has significantly improved the low energy measurements, obtaining an estimate which is about a factor of 2 lower than previous determinations."
+ Due to its significant impact on the efficiency of the CNO cycle. hence on the age - luminosity calibration during the late MS evolution. Imbriant et al. (," Due to its significant impact on the efficiency of the CNO cycle, hence on the age - luminosity calibration during the late MS evolution, Imbriani et al. ("
+2004) studied the effect of the new reaction rate on the age dating of galactic Globular Clusters (GCs).,2004) studied the effect of the new reaction rate on the age dating of galactic Globular Clusters (GCs).
+ Their analysis has shown that the new rate for the Λερ.y)O reaction leads to isochrones with a brighter and hotter Turn Off (TO) for a fixed age. and the resulting age - TO luminosity calibration predicts systematically older GC ages. by «0.9 Gyr on average.," Their analysis has shown that the new rate for the $^{14}N(p,\gamma)^{15}O$ reaction leads to isochrones with a brighter and hotter Turn Off (TO) for a fixed age, and the resulting age - TO luminosity calibration predicts systematically older GC ages, by $\sim$ 0.9 Gyr on average."
+ Weiss et al. (, Weiss et al. (
+2005) have investigated the effect of this new reaction rate on the evolution of both low- and intermediate-mass stars.,2005) have investigated the effect of this new reaction rate on the evolution of both low- and intermediate-mass stars.
+ They confirmed the results obtained by Imbriant et al. (, They confirmed the results obtained by Imbriani et al. (
+2004) and extended their analysis to more advanced evolutionary stages such as the He-burning ignition at the Red Giant Branch (RGB) tip in low-mass stars. and the core and shell He-burning stages in intermediate mass stars.,"2004) and extended their analysis to more advanced evolutionary stages such as the He-burning ignition at the Red Giant Branch (RGB) tip in low-mass stars, and the core and shell He-burning stages in intermediate mass stars."
+ More recently. Magic et al. (," More recently, Magic et al. ("
+"2010) have highlighted the crucial role played by the ""Np.y)O reaction rate in determining the transition mass to stars harbouring à convective core.","2010) have highlighted the crucial role played by the $^{14}N(p,\gamma)^{15}O$ reaction rate in determining the transition mass to stars harbouring a convective core."
+ These studies. however. do not explore several key evolutionary. properties. like the RGB bump brightness and evolutionary lifetimes during the core He-burning phase of low-mass. metal-poor stellar models. or the temperature location of models along the Horizontal Branch (HB).," These studies, however, do not explore several key evolutionary properties, like the RGB bump brightness and evolutionary lifetimes during the core He-burning phase of low-mass, metal-poor stellar models, or the temperature location of models along the Horizontal Branch (HB)."
+ The main aim of this paper is to fill this gap., The main aim of this paper is to fill this gap.
+ We analyze in detail several evolutionary properties of stellar models representative of stars currently evolving along the RGB and the HB of Galactic GCs., We analyze in detail several evolutionary properties of stellar models representative of stars currently evolving along the RGB and the HB of Galactic GCs.
+" We will also investigate the impact of the new '""""N(p.y)?O reaction rate on one of the most important primary distance indicators. te. the I-band absolute magnitude of the tip of the RGB (TRGB). and on the R- commonly used for estimating the initial Helium abundance in Galactic GCs,"," We will also investigate the impact of the new $^{14}N(p,\gamma)^{15}O$ reaction rate on one of the most important primary distance indicators, i.e. the I-band absolute magnitude of the tip of the RGB (TRGB), and on the R-parameter, commonly used for estimating the initial Helium abundance in Galactic GCs."
+ The plan of this paper is as follows: in the next section we describe briefly the set of evolutionary models employed in our analysis: Section 3 presents the impact of the new reaction rate on selected evolutionary features along the RGB stage. and provides an updated calibration of the TRGB brightness as standard candle: Section 4+ will deal with the core," The plan of this paper is as follows: in the next section we describe briefly the set of evolutionary models employed in our analysis; Section 3 presents the impact of the new reaction rate on selected evolutionary features along the RGB stage, and provides an updated calibration of the TRGB brightness as standard candle; Section 4 will deal with the core"
+representation of the C70.D) Euclidean Clifford algebra generated by Ἐν.,"representation of the $Cl(0,D)$ Euclidean Clifford algebra generated by $\Gamma_\mu$."
+ We Given D. the minimal value d. corresponding to (he ineducible representation of the Clifford algebra. is unambiguously fixed and (he maximal value ρω of the supersvimnietric extension can be computed.," We Given $D$, the minimal value $d$, corresponding to the irreducible representation of the Clifford algebra, is unambiguously fixed and the maximal value $N_{max}$ of the supersymmetric extension can be computed."
+ Expressing. lor D>1. withr—0.1.2.....T. we have. for the minimal d. and. for ων. One should note the appearance of the BRadon-Iburwitz We present. up to D«x 16. the table," Expressing, for $D\geq 1$, with$r=0,1,2,\ldots, 7$, we have, for the minimal $d$, and, for $N_{max}$, One should note the appearance of the Radon-Hurwitz We present, up to $D\leq 16$ , the table"
+Debris disks are often described as more massive versions of the Solar System's Kuiper Delt (e.g.Greavesetal.2004;Bryden2006;Jewitt2009;Booth2009).,"Debris disks are often described as more massive versions of the Solar System's Kuiper Belt \citep[e.g.][]{grea04, bryd06, jewi09, boot09}."
+". Debris disks, like the disks around Fomalhaut, Vega, e Eridani, etc."," Debris disks, like the disks around Fomalhaut, Vega, $\epsilon$ Eridani, etc."
+ can only be imaged via existing techniques if they have optical depths of —107 or higher., can only be imaged via existing techniques if they have optical depths of $\sim 10^{-4}$ or higher.
+ Models of Kuiper Belt (KB) dust production informed by dust detectors in the outer Solar System suggest a face-on optical depth of more like 107* for the KB (Backmanetal.1995;Mukai 1998).," Models of Kuiper Belt (KB) dust production informed by dust detectors in the outer Solar System suggest a face-on optical depth of more like $10^{-7}$ for the KB \citep{back95, ster96, yama98}."
+". But perhaps when the KB was younger and more massive, it closely resembled the debris disks we have seen so far around other stars."," But perhaps when the KB was younger and more massive, it closely resembled the debris disks we have seen so far around other stars."
+ This analogy has many ramifications., This analogy has many ramifications.
+" For example, images of debris disks around nearby stars show rings, clumps, warps and other asymmetries; these asymmetries have often been compared to the asymmetries in the Kuiper Belt caused by dynamical perturbations from Neptune and other planets."," For example, images of debris disks around nearby stars show rings, clumps, warps and other asymmetries; these asymmetries have often been compared to the asymmetries in the Kuiper Belt caused by dynamical perturbations from Neptune and other planets."
+" When we see these patterns in debris disks, can we recognize the planets that are sculpting them?"," When we see these patterns in debris disks, can we recognize the planets that are sculpting them?"
+" Can we use the patterns to find hidden planets that we couldn't otherwise detect, or measure the orbital parameters of planets orbiting too slowly to track?"," Can we use the patterns to find hidden planets that we couldn't otherwise detect, or measure the orbital parameters of planets orbiting too slowly to track?"
+" The Kuiper Belt, because of its proximity to the Earth, is potentially an important laboratory for testing our dynamical models of debris disks and our ideas about debris disk morphologies."," The Kuiper Belt, because of its proximity to the Earth, is potentially an important laboratory for testing our dynamical models of debris disks and our ideas about debris disk morphologies."
+ Several authors have made dynamical models of the distribution of dust in the Kuiper Belt for comparison with images of other debris disks., Several authors have made dynamical models of the distribution of dust in the Kuiper Belt for comparison with images of other debris disks.
+" Liou&Zook(1999) showed that Neptune may temporarily trap dust in mean motion resonances (MMRs), forming a wide circumsolar ring, from 35-50 AU, with a gap in the ring at the location of Neptune."," \citet{lz99} showed that Neptune may temporarily trap dust in mean motion resonances (MMRs), forming a wide circumsolar ring, from 35–50 AU, with a gap in the ring at the location of Neptune."
+" This model has often been compared to the wide, clumpy rings seen around Epsilon Eridani and Vega (e.g.Macintoshetal.2003).. Moro-Martin&M"," This model has often been compared to the wide, clumpy rings seen around Epsilon Eridani and Vega \citep[e.g.][]{maci03}."
+"alhotra(2002) explored how grains of various sizes behave in the outer Solar System, and predicted the spectral energy distribution of the Kuiper belt dust (seealsoMoro-Martin&Malhotra2003)."," \citet{mm1} explored how grains of various sizes behave in the outer Solar System, and predicted the spectral energy distribution of the Kuiper belt dust \citep[see also][]{mm2}."
+". Holmesetal.(2003) explored how a particular family of KBOs, the plutinos, could contribute to the resonant Kuiper Belt dust population."," \citet{holm03} explored how a particular family of KBOs, the plutinos, could contribute to the resonant Kuiper Belt dust population."
+ But these models contain an important limitation: they largely neglect grain-grain collisions., But these models contain an important limitation: they largely neglect grain-grain collisions.
+" For some grain sizes in any debris disk, the typical collision time becomes shorter than the typical Poynting-Robertson (PR) time, affecting the disk morphology"," For some grain sizes in any debris disk, the typical collision time becomes shorter than the typical Poynting-Robertson (PR) time, affecting the disk morphology"
+In this section. we follow the evolution. of jetted CRB remnants numerically to see what elfects will those parameters (such as £0.62.40.Usnop. sss ele.),"In this section, we follow the evolution of jetted GRB remnants numerically to see what effects will those parameters (such as $\xi_{\rm e}, \xi_{\rm B}^2, \theta_0, \theta_{\rm obs}, n, p,$ ..., etc.)"
+ have on the optical light curves., have on the optical light curves.
+" For convenience. let us define the Following initial values or parameters as a set of “stanclarel™ parameters: initial encrey per solid angle £y/Q=107! ergs/dx. y=300 (i.e. initial ejecta mass per solid. angle Aly/Qo2O.OOLQAL. /4x). 0=lem *. £3=0.01. p=2.5. Dy,=10)10"" κρο £.=0.1. 6,=0.2.4.0. where the observing angle 6,1. is defined as the angle between the line of sight and the jet axis."," For convenience, let us define the following initial values or parameters as a set of “standard” parameters: initial energy per solid angle $E_0 / \Omega_0 = 10^{54}$ $4 \pi$, $\gamma_0 = 300$ (i.e., initial ejecta mass per solid angle $M_{\rm ej}/\Omega_0 \approx 0.0019 M_{\odot}/ 4 \pi$ ), $n = 1$ $^{-3}$, $\xi_{\rm B}^2 = 0.01$, $p = 2.5$, $D_{\rm L} = 1.0 \times 10^6$ kpc, $\xi_{\rm e} =0.1$, $\theta_0 = 0.2, 
+\theta_{\rm obs}=0$, where the observing angle $\theta_{\rm obs}$ is defined as the angle between the line of sight and the jet axis."
+" For simplicity. we first assume that the expansion is completely aciabatic all the time (i.e. €—0. we call it an ""ideal jet. distinguishing it from the “realistic” jet defined in Section 2.1)."," For simplicity, we first assume that the expansion is completely adiabatic all the time (i.e. $\epsilon \equiv 0$, we call it an “ideal” jet, distinguishing it from the “realistic” jet defined in Section 2.1)."
+ Figure | shows the evolution of the Lorentz factor for some exemplary jets., Figure 1 shows the evolution of the Lorentz factor for some exemplary jets.
+" In the ""standard case. the relativistic phase lasts onlv for ~10"" s. and the non- phase begins at about /—LO! s. In short. the ejecta. will cease to be highly relativistic at time /107 Q( s. OPENThis again. gives. strong support to our previous. argument that we should be careful in discussing the fireball evolution under the simple assumption of ultra-relativistic limit (Lluane et al."," In the “standard” case, the ultra-relativistic phase lasts only for $\sim 10^5$ s, and the non-relativistic phase begins at about $t \sim 10^{6.5}$ s. In short, the ejecta will cease to be highly relativistic at time $t \sim 10^5$ — $10^6$ s. This again gives strong support to our previous argument that we should be careful in discussing the fireball evolution under the simple assumption of ultra-relativistic limit (Huang et al."
+ 1998a. b. 1999a. b).," 1998a, b, 1999a, b)."
+" In the case of a dense ISM (yp=10"" 7. the clash-dotted lino). the expansion will become non-relativistic as carly as (£710077 s. Figure 2 js the evolution of the jet opening angle 6."," In the case of a dense ISM $n = 10^6$ $^{-3}$, the dash-dotted line), the expansion will become non-relativistic as early as $t \sim 10^{4.5}$ s. Figure 2 is the evolution of the jet opening angle $\theta$."
+ During the ultra-relativistic phase. 6 increases only slightly.," During the ultra-relativistic phase, $\theta$ increases only slightly."
+ But at the Newtonian stage. the increase of 6 is very quick.," But at the Newtonian stage, the increase of $\theta$ is very quick."
+" The effect. of €& on the optical (It-band) light curves is illastratec in Figure 3. from which we can see clearly that: (1) In no case could we observe the theoretically predicted light ""urve steepenine (with the break point determined by 5~ 4) duringself. this is consistent with the result of Alocderski. Sikora Bulik (1999)."," The effect of $\xi_{\rm e}$ on the optical (R-band) light curves is illustrated in Figure 3, from which we can see clearly that: (i) In no case could we observe the theoretically predicted light curve steepening (with the break point determined by $\gamma \sim 
+1/\theta$ ) during, this is consistent with the result of Moderski, Sikora Bulik (1999)."
+ Note that in his ligure the relativistic phase is restricted by f<LO? s. lluangὃν et al. (, Note that in this figure the relativistic phase is restricted by $t \leq 10^6$ s. Huang et al. (
+1999¢) have ogiven a reasonable explanation to his phenomena: the edge and the lateral expansion ellects xin to take elfect when 5~1/6 (in our caleulation this occurs at /1077 s. however the blastwave is already in its milclly relativistic phase at that moment and it will become non-relativistic soon after that (le... when /2LO! s. see Figure 1).,"1999c) have given a reasonable explanation to this phenomena: the edge and the lateral expansion effects begin to take effect when $\gamma \sim 1/\theta$ (in our calculation this occurs at $t \sim 10^{5.5}$ s, however the blastwave is already in its mildly relativistic phase at that moment and it will become non-relativistic soon after that (i.e., when $t \geq 10^{6.5}$ s, see Figure 1)."
+ So it is not surprising that we could not see any obvious breaks during the relativistic phase. they just do not have time to emerge.," So it is not surprising that we could not see any obvious breaks during the relativistic phase, they just do not have time to emerge."
+ Another possible reason is: at. these stages. since 5 is no longer much larger than 1. we should be careful in using the analytic approximations: (ii) When £. is small. an obvious break does appear in the light curve. but," Another possible reason is: at these stages, since $\gamma$ is no longer much larger than 1, we should be careful in using the analytic approximations; (ii) When $\xi_{\rm e}$ is small, an obvious break does appear in the light curve, but"
+"to the jet (henceforth called ECS models for ""external Compton scattering).",to the jet (henceforth called ECS models for “external Compton scattering”).
+ schlickeiser(1993) use the accretion disk as (he source of soft photons whereas others (Ghisellini&Macau1996) use broac-line region clouds as the source of soft photons.," \citet{der93} use the accretion disk as the source of soft photons whereas others \citep{ghi96}
+ use broad-line region clouds as the source of soft photons."
+ eince the end of the EGRET mission there have also been several reanalvses of the significance of iclentifications of EGRET sources. particularly those of Mattoxefal.(1991):Mattox.Hartman&Reimer(2001) ancl Sowarcds-Enunerd.Romani.&Michelson(2003.," Since the end of the EGRET mission there have also been several reanalyses of the significance of identifications of EGRET sources, particularly those of \citet{mat97,mat01} and \citet{sow03,sow04}."
+ 2004)..The Sowards-Emmerd.Romani.&Michelson(2003.2004) survey exclucles sources with |b}<10 (the Mattox papers exclude sources with |b]<3°). and thus may exclude additional sources that are thought to be blazars (SeneraefaL.2004).," .The \citet{sow03,sow04} survey excludes sources with $|b|<10^{\circ}$ (the Mattox papers exclude sources with $|b|<3^{\circ}$ ), and thus may exclude additional sources that are thought to be blazars \citep{sgu04}."
+. Likewise. some work continues on identilvine other individual sources. such as 3EG JO416+3650. possibly identified with 3C 111 (SeneraefaL.2006).," Likewise, some work continues on identifying other individual sources, such as 3EG J0416+3650, possibly identified with 3C 111 \citep{sgu06}."
+. In the eight vears since the end of the CGRO mission. (here has been substantial modification to the identifications eiven in (he last catalog published by the EGRET team Clartinanefaf1999).. so we present a comprehensive statistical analvsis including all potential identifications. using homogeneous criteria [or inclusion in (he sample (82).," In the eight years since the end of the CGRO mission, there has been substantial modification to the identifications given in the last catalog published by the EGRET team \citep{har99}, so we present a comprehensive statistical analysis including all potential identifications, using homogeneous criteria for inclusion in the sample )."
+ We have also included various statistical techniques of survival analvsis for the inclusion of upper aud lower limits in the data (83)., We have also included various statistical techniques of survival analysis for the inclusion of upper and lower limits in the data ).
+ We later discuss whether or not the results of the analvses are dependent on the precise source list., We later discuss whether or not the results of the analyses are dependent on the precise source list.
+ The approach used in understanding the physical implications of our results is to use a Monte Carlo technique (Lister&Marscher1997) to generate multiple simulated samples under differing assumptions of overall theoretical or phenomenological model (84.5). distribution of physical parameters such as bulk Lorentz lactor (e.g.. Gaussian. power-law. ete.)," The approach used in understanding the physical implications of our results is to use a Monte Carlo technique \citep{lis97} to generate multiple simulated samples under differing assumptions of overall theoretical or phenomenological model ), distribution of physical parameters such as bulk Lorentz factor (e.g., Gaussian, power-law, etc.)"
+ and selection effects., and selection effects.
+ Adcditional physical implications of observed correlations and implications lor the Gamma Rav Large Area Space Telescope (GLAST). successfully launched in June 2008. are discussed in86.," Additional physical implications of observed correlations and implications for the Gamma Ray Large Area Space Telescope (GLAST), successfully launched in June 2008, are discussed in."
+7. To construct a comprehensive list of identified eamuima-ray blazars we have started with the Third EGRET Catalog (Hartmanefa£.1999)., To construct a comprehensive list of identified gamma-ray blazars we have started with the Third EGRET Catalog \citep{har99}.
+. We note thal we are not re-evaluating the detections themselves., We note that we are not re-evaluating the detections themselves.
+ To add or remove sources [rom this list of identifications using uniform criteria. we have further used the guidance of Sowards-Emmered.Romani.son(2003. 2004).," To add or remove sources from this list of identifications using uniform criteria, we have further used the guidance of \citet{sow03,sow04}."
+. Though other standards lor inclusion or rejection do exist. we have used these particular survevs to keep inclusion criterion standard.," Though other standards for inclusion or rejection do exist, we have used these particular surveys to keep inclusion criterion standard."
+ We have taken the best blazar identifications with Figure of Merit (FoM) criterion greater than 0.25 (Sowards-Emmerd.mani.&Michelson2003. 2004).," We have taken the best blazar identifications with Figure of Merit (FoM) criterion greater than 0.25 \citep{sow03,sow04}."
+. The FoM used here is a statistic that incorporates both radio and X-ray spectral information in evaluating the probability that à particular radio source is a malch to the gamma-ray source., The FoM used here is a statistic that incorporates both radio and X-ray spectral information in evaluating the probability that a particular radio source is a match to the gamma-ray source.
+ For sources with multiple possible identifications. we include," For sources with multiple possible identifications, we include"
+a microlensed circular source. which peaks at.,"a microlensed circular source, which peaks at."
+1004... Phe same feature is seen in all lisht. curves. namely that the observed. [lux drops as the planetary phase. changes fron full to crescent. even though these latter phases undergo stronger magnification.," The same feature is seen in all light curves, namely that the observed flux drops as the planetary phase changes from full to crescent, even though these latter phases undergo stronger magnification."
+ Fieure 7 presents simulated. observations. of light curves given in Figures 2 and 4. assuming the host star is V—21S. with an extinction. of L25maes (following the model of LI2000).," Figure \ref{noise} presents simulated observations of light curves given in Figures 2 and 4, assuming the host star is V=18, with an extinction of 1.25mags (following the model of LI2000)."
+ The. three models. (a). (b) ancl (c) are. presented vertically. with each column. representing he dillering orientations.," The three models, (a), (b) and (c) are presented vertically, with each column representing the differing orientations."
+ For each model. two sets of simulated. observations are given. assuming a l10-m class (upper panels) ancl 30-m2 class (lower panels). using the xdanetary. parameters used ον LI2000.," For each model, two sets of simulated observations are given, assuming a 10-m class (upper panels) and 30-m class (lower panels), using the planetary parameters used by LI2000."
+ The ordinate on each yanel represents. the fractional light curve deviation. due o the presence of the planet. and. should. be compared to Figure 2 in 1192000.," The ordinate on each panel represents the fractional light curve deviation due to the presence of the planet, and should be compared to Figure 2 in LI2000."
+ Each telescope is assumed to have a total elliciency.. anc the data points are. subject to whoton counting errors only.," Each telescope is assumed to have a total efficiency, and the data points are subject to photon counting errors only."
+ The integration times for both elescopes are. LOO. 200 and 400secs for model (a). (b) and (c) respectively.," The integration times for both telescopes are 100, 200 and 400secs for model (a), (b) and (c) respectively."
+ The I0-mi class telescope observations represent the current astronomical technology., The 10-m class telescope observations represent the current astronomical technology.
+ The fractional deviation introciuced into a light curve by model (a). of order ~254. is ‘learly detectable. with the photometry accurately revealing 1e dillering light curve profiles at the cdillerent. planetary srlentations.," The fractional deviation introduced into a light curve by model (a), of order $\sim2\%$, is clearly detectable, with the photometry accurately revealing the differing light curve profiles at the different planetary orientations."
+ Ες is also seen with moclel (b) which. while possessing lower signal-to-noise. still allows recovery of the etailed light curve profile.," This is also seen with model (b) which, while possessing lower signal-to-noise, still allows recovery of the detailed light curve profile."
+ For model (0). however. the small fraction of rellected lisht means that the resultant fractional deviations are small. less than4.," For model (c), however, the small fraction of reflected light means that the resultant fractional deviations are small, less than."
+. As illustrated. in ligure 7.. while Hluctuations at this level are potentially etectable. the details of their profiles are lost in the noise.," As illustrated in Figure \ref{noise}, while fluctuations at this level are potentially detectable, the details of their profiles are lost in the noise."
+ Plans are currently. underway to build the next. generation of ground-based. telescope. with 302m. telescopes entering," Plans are currently underway to build the next generation of ground-based telescope, with 30-m telescopes entering"
+path 1.,path 1.
+ After the binary has moved the distance £ from the centre of the cluster. O. it clisintegrates at the position 7? and one component of the binary is released.," After the binary has moved the distance $\xi$ from the centre of the cluster, $O$, it disintegrates at the position $P$ and one component of the binary is released."
+ The star cluster. from which the star has been ejected. is identified. if the extrapolated path of the released star intersects a sphere with the identification radius. #2). round the star cluster.," The star cluster, from which the star has been ejected, is identified, if the extrapolated path of the released star intersects a sphere with the identification radius, $R_\mathrm{i}$, round the star cluster."
+ If the released star moves along path 2. the star cluster is identified. (extrapolated long dashed line).," If the released star moves along path 2, the star cluster is identified (extrapolated long dashed line)."
+ Lf the star moves along path 3. the parent star cluster can not been identified (extrapolated short dashed line).," If the star moves along path 3, the parent star cluster can not been identified (extrapolated short dashed line)."
+ Only stars which have new orbits within the grav shaded region between the dotted lines can be traced back to their parent star cluster., Only stars which have new orbits within the gray shaded region between the dotted lines can be traced back to their parent star cluster.
+ At the location of binary disintegration (point 7?) the identification sphere appears under the angle 2a. ie. a solid angle of 2r(lcosa)," At the location of binary disintegration (point $P$ ) the identification sphere appears under the angle $2 \alpha$ , i.e. a solid angle of $2 \pi (1-\cos\alpha)$."
+ Because the directions of the released stars are rancomly distributed over the full solid angle of dr. the probability that the star cluster can be identified is given by the ratio of the solid angele of the identification sphere ancl the full solid angle. The cosine can be expressed by the identification radius. I. and the disintegrationD distance. M£. Then the identification probability is The resulting probabilities for dillerent. ratios of the identification racius ancl disintegration distance are listed in Table 1..," Because the directions of the released stars are randomly distributed over the full solid angle of $4 \pi$, the probability that the star cluster can be identified is given by the ratio of the solid angle of the identification sphere and the full solid angle, The cosine can be expressed by the identification radius, $R_\mathrm{i}$, and the disintegration distance, $\xi$, Then the identification probability is The resulting probabilities for different ratios of the identification radius and disintegration distance are listed in Table \ref{tab_probability}."
+ Alassive stars ancl massive binaries can only be ejected cdvnamically from star clusters during the early stage of their life., Massive stars and massive binaries can only be ejected dynamically from star clusters during the early stage of their life.
+ Lf they form within compact few-body configurations or trapezia systems. (2).. then the decay. time scale of these few-hocly systems (1 Myr. 2)) implies only early ejections of massive stars.," If they form within compact few-body configurations or trapezia systems \citep{clarke1992a}, , then the decay time scale of these few-body systems $<$ 1 Myr, \citealp{pflamm-altenburg2006a}) ) implies only early ejections of massive stars."
+ Lf massive stars ancl binaries are formed distributed. over the star cluster. then ejections can only occur as long as the stellar densities are high.," If massive stars and binaries are formed distributed over the star cluster, then ejections can only occur as long as the stellar densities are high."
+ Due to gas expulsion voung embedded star clusters become super-virial ancl start to expand and the stellar density decreases rapidly (?).., Due to gas expulsion young embedded star clusters become super-virial and start to expand and the stellar density decreases rapidly \citep*{kroupa2001b}.
+ The time scale of decrease. of the stellar density is comparable to the eas expulsion time scale. of the order of σε) Myr.," The time scale of decrease of the stellar density is comparable to the gas expulsion time scale, of the order of $\approx$ 1 Myr."
+ Thus the time of flight of ejected. binaries is comparable to their maximum-elifetimo., Thus the time of flight of ejected binaries is comparable to their maximum-lifetime.
+ ‘Taking a lower velocity cut-olf for O-star runaways of 30 knis tas considered in ? and a mean life-time of 5 Myr of O-stars implies a clisintegration distance £ —150 pe., Taking a lower velocity cut-off for O-star runaways of 30 km $^{-1}$ as considered in \citet{gies1986a} and a mean life-time of 5 Myr of O-stars implies a disintegration distance $\xi=$ 150 pc.
+ For a large identification radius of //; 210 pe of the star cluster. the back-tracing probability is 1 ραδιο 1)).," For a large identification radius of $R_\mathrm{i}=$ 10 pc of the star cluster, the back-tracing probability is 1 (Table \ref{tab_probability}) )."
+ Lowering the cut-olf velocity of the runaway-star definition to 10 kim s results in a back-tracing probability of 1 per cent., Lowering the cut-off velocity of the runaway-star definition to 10 km $^{-1}$ results in a back-tracing probability of 1 per cent.
+ Thus. ected massive binaries which are listed in runaway O-star surveys will produce O stars which can not be traced back to their parent star cluster.," Thus, ejected massive binaries which are listed in runaway O-star surveys will produce O stars which can not be traced back to their parent star cluster."
+ But one might expect that in this case the massivestar can be traced back to à supernova shell., But one might expect that in this case the massivestar can be traced back to a supernova shell.
+ However. as single supernova shells disappear on a time scale of 0.51 Myr (2). it might be possible to identify the parent supernova of the released star only in very rare Cases.," However, as single supernova shells disappear on a time scale of 0.5–1 Myr \citep{chevalier1974a} it might be possible to identify the parent supernova of the released star only in very rare cases."
+ The observationally derived: runaway fraction among O stars varies. widely in the literature. (see. for example Table 19 in ? or Table 5 in ?2)).," The observationally derived runaway fraction among O stars varies widely in the literature, (see for example Table 13 in \citealp{gies1986a} or Table 5 in \citealp{stone1991a}) )."
+ 7. identified 15 starsout of a sample of 36 runaway candidates with confirmed peculiar radial velocities 2230 kms |., \citet{gies1986a} identified 15 starsout of a sample of 36 runaway candidates with confirmed peculiar radial velocities $\ge$ 30km $^{-1}$ .
+ Comparing with their, Comparing with their
+The survey volume for this analysis is similar to that described iu refsec-voL.,The survey volume for this analysis is similar to that described in \\ref{sec-vol}.
+ An important difference is that we place a stricter eut on the point spread finetion (see relsec-data)) aud iu the process lose many de Vaucouleurs galaxies., An important difference is that we place a stricter cut on the point spread function (see \\ref{sec-data}) ) and in the process lose many de Vaucouleurs galaxies.
+ In it was noted that there is a Iuminosity — axis ratio trend., In \citet{vr05} it was noted that there is a luminosity – axis ratio trend.
+ In the range [718.220] the trend is approximately flat.," In the range $[-18, -20]$ the trend is approximately flat."
+ In the range [-20.—22] the axis ratio increases with increasing luminosity.," In the range $[-20, -22]$ the axis ratio increases with increasing luminosity."
+ For galaxies brighter than —22. the axis ratio seems to level off.," For galaxies brighter than $-22$, the axis ratio seems to level off."
+ Iu principle. this correlation could affect our results.," In principle, this correlation could affect our results."
+ This motivates us to keep luminosity bius as narrow as possible. so as to ninunize any effect of this trend within a single bin.," This motivates us to keep luminosity bins as narrow as possible, so as to minimize any effect of this trend within a single bin."
+ Ou the other haud. statistical claims are ouly useful lor large saliple sizes.," On the other hand, statistical claims are only useful for large sample sizes."
+ To keep our subsatmples of galaxies large enough. we are motivated to keep luminosity bins as wide as possible.," To keep our subsamples of galaxies large enough, we are motivated to keep luminosity bins as wide as possible."
+ As a Compromise. we make Duminosity bins two magnitudes wide.," As a compromise, we make luminosity bins two magnitudes wide."
+ We start by creating three volume-limited surveys with absolute magnitude cut-offs —]18.—20. aud —22.," We start by creating three volume-limited surveys with absolute magnitude cut-offs $M_{r_0} = -18, -20,$ and $-22$."
+" For the volumes correspoucing to the first two cut-olls. we cletermine the euvironiuent (see 'elsec-env)) for target. galaxies that have absolute magnitude Al,&[M,,.M,—2]. while for the wightest cut-off. we find the environment for all galaxies with Al,€[-22.—»c]."," For the volumes corresponding to the first two cut-offs, we determine the environment (see \\ref{sec-env}) ) for target galaxies that have absolute magnitude $M_r \in [M_{r_0}, M_{r_0}-2]$, while for the brightest cut-off we find the environment for all galaxies with $M_r \in [-22, -\infty]$."
+ There are only sixteen galaxies that are brighter than —21 in absolute maguituce in our volume: therefore. their inclusion does uot significantly widen the effective magnitude rauge of the biu.," There are only sixteen galaxies that are brighter than $-24$ in absolute magnitude in our volume; therefore, their inclusion does not significantly widen the effective magnitude range of the bin."
+" For example. in the volume corresponding to the A,=—18 cut-off we find the environment for all galaxies with where the luminosity distauce ofa galaxy is no larger than dy,= 113 Mpe. correspoucing o the riaximum luminosity distance at which a galaxy of absolute magnitude A4,=—18 could be observed."," For example, in the volume corresponding to the $M_{r_0}=-18$ cut-off we find the environment for all galaxies with $ M_r \in [-18,-20]$ where the luminosity distance of a galaxy is no larger than $d_L=$ 143 Mpc, corresponding to the maximum luminosity distance at which a galaxy of absolute magnitude $M_{r_0}=-18$ could be observed."
+ The total παο of galaxies in each volume. aud the percentage of each profile type is oesented in Table 2..," The total number of galaxies in each volume, and the percentage of each profile type is presented in Table \ref{tab:shape}."
+ Table 2. also lists. under the headiug evlinder width. the total number of cylinders (see 'e[sec-euv)) that do not intersect the border of the SDSS spectrographie survey. which is the same as the number of target. galaxies that we keep in our aualvsis.," Table \ref{tab:shape} also lists, under the heading `cylinder width', the total number of cylinders (see \\ref{sec-env}) ) that do not intersect the border of the SDSS spectrographic survey, which is the same as the number of target galaxies that we keep in our analysis."
+ The galaxies iu each sample are subject to quality cuts which are described iu relsec-data.., The galaxies in each sample are subject to quality cuts which are described in \\ref{sec-data}.
+ The binuine procedure is described iu refsec-env.., The binning procedure is described in \\ref{sec-env}.
+ Figures 5 through 7 show the correlations of apparent axis ratio with environment. labeled by profile type.," Figures \ref{fig:5} through \ref{fig:7} show the correlations of apparent axis ratio with environment, labeled by profile type."
+ Tables 3. through 5 contain all related. statistical information., Tables \ref{tab:shape18} through \ref{tab:shape22} contain all related statistical information.
+ We plot a best fit straight line through the binned data that minimizes the chi-square., We plot a best fit straight line through the binned data that minimizes the chi-square.
+ The value of tlie reduced chi-square for each fit is printed in the tables., The value of the reduced chi-square for each fit is printed in the tables.
+ The error bars show the error in the mean assumiug, The error bars show the error in the mean assuming
+]t is known that catastrophic changes in (he magnetic field configuration of active regions (ARs) at coronal heights trigger energetic explosive events such as flares and coronal mass ejections (CMES).,It is known that catastrophic changes in the magnetic field configuration of active regions (ARs) at coronal heights trigger energetic explosive events such as flares and coronal mass ejections (CMEs).
+ Therefore. variations in magnetic and Doppler velocity fields are expected lo accompany these energetic events.," Therefore, variations in magnetic and Doppler velocity fields are expected to accompany these energetic events."
+ It was first suggested by Giovanelli(1939) thal energy release in flares should be associated with observable magnetic field changes., It was first suggested by \citet{Giovanelli1939} that energy release in flares should be associated with observable magnetic field changes.
+ Extensive, Extensive
+(8) The power-law indices zz4.4—7 estimated directly. [rom the observed distributions differ ereatlv from the value zz1 for simple SOC theory aud lie outside (he range +0.5—3 of known SOC systems.,(8) The power-law indices $\approx 4.4 - 7$ estimated directly from the observed distributions differ greatly from the value $\approx 1$ for simple SOC theory and lie outside the range $\approx 0.5 - 3$ of known SOC systems.
+ SOC is therefore implausible despite the statistics being power-, SOC is therefore implausible despite the statistics being power-law.
+ Reversing (his conclusion requires. first. Further observations confirming that (he giant phenomena have intrinsic power-law statistics whose indices were signilicantly overestimated due to convolution effects with the normal pulses ancl measurement noise and/or. second. construction of tlieoretical SOC models with suitably large values of a. (," Reversing this conclusion requires, first, further observations confirming that the giant phenomena have intrinsic power-law statistics whose indices were significantly overestimated due to convolution effects with the normal pulses and measurement noise and/or, second, construction of theoretical SOC models with suitably large values of $\alpha$. ("
+9) Nonlinear modulational instabilities and wave collapse processes (vpically have power-law field statistics. qualitatively consistent with the available data.,"9) Nonlinear modulational instabilities and wave collapse processes typically have power-law field statistics, qualitatively consistent with the available data."
+ Moreover. the observed indices azz4.4—7 lie within the range 4-7 predicted by current theories for wave collapse (Robinson&Newman1990:Robinson1996).," Moreover, the observed indices $\alpha \approx 4.4 - 7$ lie within the range $4 - 7$ predicted by current theories for wave collapse \citep{rn1990,r1996}."
+. This is currently (the most plausible mechanism lor giant pulses and micropulses. complementing another argument based on timescales (ILankinsetal.M2003).. (," This is currently the most plausible mechanism for giant pulses and micropulses, complementing another argument based on timescales \citep{hankinsetal2003}. ("
+10) Further research on applying collapse theory ancl modulational instabilities to giant pulses and micropulses is required (o resolve open issues relating to (1) whether the field statistics for collapse in the relativistic. electron-positron. highly magnetized plasmas in pulsar magnetospheres differ Irom those known lor non-relativistic electron-proton plasmas. Gi) explaining the apparent (rend in a with observing frequency [or giant pulses being opposite to that predicted «qualitatively based on magnetization effects and current collapse theory. and (iii) considering in detail how collapse produces the observed radiation aud how anv ancillary radiation mechanism mocdifies (he predicted field statistics. (,"10) Further research on applying collapse theory and modulational instabilities to giant pulses and micropulses is required to resolve open issues relating to (i) whether the field statistics for collapse in the relativistic, electron-positron, highly magnetized plasmas in pulsar magnetospheres differ from those known for non-relativistic electron-proton plasmas, (ii) explaining the apparent trend in $\alpha$ with observing frequency for giant pulses being opposite to that predicted qualitatively based on magnetization effects and current collapse theory, and (iii) considering in detail how collapse produces the observed radiation and how any ancillary radiation mechanism modifies the predicted field statistics. ("
+11) Finally. SGT cannot vel be ruled out as important in understanding giant. pulses and micropulses since (i) the observed power-law statistics can be interpreted in terms of vector convolution of one or more of the lognormal distributions predicted lor SGT. as already observed for normal pulses (Cairnsetal.2002.2003b.c).. and Gi) wave collapse and SGT can coexist.,"11) Finally, SGT cannot yet be ruled out as important in understanding giant pulses and micropulses since (i) the observed power-law statistics can be interpreted in terms of vector convolution of one or more of the lognormal distributions predicted for SGT, as already observed for normal pulses \citep{cetal2002d,cetal2002b,cetal2002c}, and (ii) wave collapse and SGT can coexist."
+ The differences between normal and giant pulsar emission favor. but do not require. a different theoretical interpretation for the (wo phenomena. (hereby favoring wave collapse over SGT for giant pulses and micropuUV.es.," The differences between normal and giant pulsar emission favor, but do not require, a different theoretical interpretation for the two phenomena, thereby favoring wave collapse over SGT for giant pulses and micropulses."
+ In conclusion. progress has been mace in determining the intrinsic fiekl statistics of ejant micropulses and giant pulses and identilving therefrom the source physics and emission mechanisms.," In conclusion, progress has been made in determining the intrinsic field statistics of giant micropulses and giant pulses and identifying therefrom the source physics and emission mechanisms."
+ Both phenomena appear lo have verv similar field statistics ancl to adimit a common theoretical interpretation. although whether the observed. power-law features are intrinsic or due (o convolution effects remains to be determined.," Both phenomena appear to have very similar field statistics and to admit a common theoretical interpretation, although whether the observed power-law features are intrinsic or due to convolution effects remains to be determined."
+ The nonlinear process of wave collapse produces power-law statistics wilh indices that can be in the observed range aud. despite some theoretical difliculties. appears the most plausible theoretical interpretation.," The nonlinear process of wave collapse produces power-law statistics with indices that can be in the observed range and, despite some theoretical difficulties, appears the most plausible theoretical interpretation."
+ SGT remains viable although less favored. while SOC. refractive lensing. and certain other," SGT remains viable although less favored, while SOC, refractive lensing, and certain other"
+fainter than ο=21.0 the separation between stars axd ealaxies starts to be probleuatic.,fainter than $i' = 21.0$ the separation between stars and galaxies starts to be problematic.
+ From Fig., From Fig.
+ l dt is clear that for the D1 field the contanunation of galaxies is stnall for /<21.0 while for the D2 aud the D3 field the star/galaxw separation stars to fail at already /—20.5," \ref{star-gal-separation} it is clear that for the D1 field the contamination of galaxies is small for $i' < 21.0$ while for the D2 and the D3 field the star/galaxy separation starts to fail at already $i'
+\sim 20.5$."
+ The cvice of the cit-off at /8o-21.0 for the Τους] release is certainly a conservative criterion which cau not ος applied for all three fields., The choice of the cut-off at $i' = 21.0$ for the T0001 release is certainly a conservative criterion which can not be applied for all three fields.
+ However. the star/ealaxy separation depends oi the colour of the objects.," However, the star/galaxy separation depends on the colour of the objects."
+ If one restricts to red objects with ilπ Lhesars Call ο better seperated from galasdes iud thus stars can be extracted until //«22.5 (for the D] field).," If one restricts to red objects with $r'-i'
+> 1.4$, stars can be better seperated from galaxies and thus stars can be extracted until $i' < 22.5$ (for the D1 field)."
+ We discuss aud quantity below galaxy coutamunalon as a function of uaeuitude aud colour., We discuss and quantify below galaxy contamination as a function of magnitude and colour.
+ Iun order to assess fje nununber of ealaxies contaminating our stellar sample. we used. spectroscopic data from the VIMIOS-VLT Deep Sirvey (VVDS).," In order to assess the number of galaxies contaminating our stellar sample, we used spectroscopic data from the VIMOS-VLT Deep Survey (VVDS)."
+ The VVDS (Le Fevyre ct al. 2 053) , The VVDS (Le Fèvvre et al. \cite{LeFevre2005}) )
+intends to CASTILE redshifts over θc:<5 across qu:16 dee? 2.in fourE separate fields., intends to measure redshifts over $0<z<5$ across 16 $^2$ in four separate fields.
+ The survey is couducte Lin two steps: an inaeiug survey and a spectroscopic strvov., The survey is conducted in two steps: an imaging survey and a spectroscopic survey.
+" Deep imaging in the UBVRI bands aud partly A"" baud is obtained with the CFIIT. ESO-NTT and ESO-2.2u1. The VLT VIAIOS iustrunents allow the measurement of redshifts of objects Xected from the imuagimeg survey."," Deep imaging in the UBVRI bands and partly $K'$ band is obtained with the CFHT, ESO-NTT and ESO-2.2m. The VLT VIMOS instruments allow the measurement of redshifts of objects selected from the imaging survey."
+" The so-called ""Deep survey has a linaΠιο magnitude Lip = 21.", The so-called “Deep” survey has a limiting magnitude $_{AB}$ = 24.
+ Its location. a= 2h 26 mand à = 30° overlaps with that of the D1 field.," Its location, $\alpha$ = 2 h 26 m and $\delta$ = $-4^\circ 30'$ overlaps with that of the D1 field."
+ For a ¢ctailed description of the VVDs data We refer to Le Fevvre et al. (2005))., For a detailed description of the VVDs data we refer to Le Fèvvre et al. \cite{LeFevre2005}) ).
+ We extended. the Iuuitiug magnitue for the star/ealaxy separation of the D1 field to //=22.0:00g (thus πας deeper than the official TOO01 reease]. using the half light radius. as descenbec above.," We extended the limiting magnitude for the star/galaxy separation of the D1 field to $i' = 22.0\,mag$ (thus mag deeper than the official T0001 release), using the half light radius, as described above."
+" The D1 stacks have better secius than the other fields. alowing star-ealaxy separation to fainter nmuaenituces,"," The D1 stacks have better seeing than the other fields, allowing star-galaxy separation to fainter magnitudes."
+ We use this deeper catalogue iu order to jscuss f1ο ealaxy contamination (see below)., We use this deeper catalogue in order to discuss the galaxy contamination (see below).
+ However. for the rest of the paper we use the official TOOOL release. oeji Ισ stars are separated from galaxies ouly uutil =21.0.," However, for the rest of the paper we use the official T0001 release, in which stars are separated from galaxies only until $i'=21.0$."
+" We cross-identified the D1 data with the VVDS data sing a search radius of2"".", We cross-identified the D1 data with the VVDS data using a search radius of.
+ Due to the masking aud ouly partial overlap between the DI field aud the VVDS field. only an area of 0.1 sq degree is iu comunon.," Due to the masking and only partial overlap between the D1 field and the VVDS F02-field, only an area of 0.4 sq degree is in common."
+ Out of 9088 sources of VVDS daa. TLLO sources have been cross-1identified with the D1 ποd.," Out of 9088 sources of VVDS data, 7110 sources have been cross-identified with the D1 field."
+ 295 sources were identified as stellar ancl are shown in Fig. 2.., 295 sources were identified as stellar and are shown in Fig. \ref{DEEP}.
+ Figure 2 shows the colour-colour diagrams of the D1 field cross-identified witi the VVDS data., Figure \ref{DEEP} shows the colour-colour diagrams of the D1 field cross-identified with the VVDS data.
+ All objects shown are classified as stars from the morphological criterion., All objects shown are classified as stars from the morphological criterion.
+ The dos show the spectroscopically identified stars aud asterisks εtose. objects spectroscopically ideutified as ealaxies., The dots show the spectroscopically identified stars and asterisks those objects spectroscopically identified as galaxies.
+ The majontv of the ealaxics classified morphologically as stars poplate the blue part of the colom-colour diagrams with 5—/<0.L., The majority of the galaxies classified morphologically as stars populate the blue part of the colour-colour diagram with $ r'-i' < 0.4$.
+ They populate the stellar locus as well as regions outside of the stellar. sequeicc., They populate the stellar locus as well as regions outside of the stellar sequence.
+ Table 2 eives the percentage of ealaxies contanunating the stellar sample as a fiction of / magnitude aud 1 colour., Table \ref{contamination} gives the percentage of galaxies contaminating the stellar sample as a function of $i'$ magnitude and $r'-i'$ colour.
+ The percentage of ealaxy contanuiation depeuds very much on he r! colour. which means that going to redder colour cdiuünishes seuificautlv the contamination by ealaxies.," The percentage of galaxy contamination depends very much on the $r'-i'$ colour, which means that going to redder colour diminishes significantly the contamination by galaxies."
+ For 1—7L5 aud 7o«22 one obtains for the D1 field a rate of of contamunatine galaxies while for 5—/>1.1 he percentage of galaxy contaniinatio ris negligible.," For $ r'-i' < 0.5$ and $
+i' < 22$ one obtains for the D1 field a rate of of contaminating galaxies while for $ r'-i' > 1.4$ the percentage of galaxy contamination is negligible."
+ Note hat the galaxy couanuuation is larger for the D2 aud D3 Ποια., Note that the galaxy contamination is larger for the D2 and D3 field.
+ Iu acldition. we visually inspectec (for the DI field) sources Which are ocated outside the stellar locus.," In addition, we visually inspected (for the D1 field) sources which are located outside the stellar locus."
+ We rave marked them 1i Fig., We have marked them in Fig.
+ 2 as open squares., \ref{DEEP} as open squares.
+ Most of these, Most of these
+According to the calculated. orbital period evolutions. LMXBDs can be classified into three categories:5 the divergingSHS svstems wilh £y]I2.1 the convergingSHS. svstems with «I1.1 and the parallel svstems with Py~2.,"According to the calculated orbital period evolutions, LMXBs can be classified into three categories: the diverging systems with $P_{\rm
+f} \gg P_{\rm i}$, the converging systems with $P_{\rm f} \ll P_{\rm i}$, and the parallel systems with $P_{\rm f} \sim P_{\rm i}$."
+ As an example. we present the ealeulated results for a 1.444.41.04. binary in model 3. to illustrate the three kinds of evolutionary sequences in Fig. 6..," As an example, we present the calculated results for a $1.4M_\sun+1.0M_\sun$ binary in model 3, to illustrate the three kinds of evolutionary sequences in Fig. \ref{tp}."
+ The corresponding bifurcation period is found to be 1.25 day. aud (he initial orbital periods are chosen to be 2=1.20. 1.25. and 1.40 days. which represent the converging. parallel. and civereing svstems respectively.," The corresponding bifurcation period is found to be $1.25$ day, and the initial orbital periods are chosen to be $P_{\rm
+i}=1.20$, $1.25$, and $1.40$ days, which represent the converging, parallel, and diverging systems respectively."
+ Pylvser&Savonije(1988) have emphasized the effect of MB on (he evolution of LMXDs., \citet{pylyser88} have emphasized the effect of MB on the evolution of LMXBs.
+ Comparing the results of models 1 and 2 presented in Table 2.. we find that the bifurcation periods with traditional MD law are smaller (larger) (han those with saturated MIB law. when the initial secondary star mass Mo; is less (arger) than 0.7M...," Comparing the results of models 1 and 2 presented in Table \ref{tab1}, we find that the bifurcation periods with traditional MB law are smaller (larger) than those with saturated MB law, when the initial secondary star mass $M_{\rm 2,i}$ is less (larger) than $0.7M_\sun$."
+ Our results suggest that mass loss also influences the value of P. though in an less important wav compared with MD.," Our results suggest that mass loss also influences the value of $P_{\rm bif}$, though in an less important way compared with MB."
+ The bifurcation periods in non-conservative models 3 and 4 are lower than those in model 2. in which conservative mass (transfer has been assumed.," The bifurcation periods in non-conservative models 3 and 4 are lower than those in model 2, in which conservative mass transfer has been assumed."
+ This result is consistent with vanderSluvsetal.(2005b).. but contradicted with al. (1998).," This result is consistent with \citet{sluys05b}, but contradicted with \citet{ergma98}."
+. It is also interesting to see whether an UCXD can form with saturated. MD., It is also interesting to see whether an UCXB can form with saturated MB.
+ For an LMXD with an initial orbital period below the bifurcation period. mass (ransfer is mainly driven by the loss of angular momentum.," For an LMXB with an initial orbital period below the bifurcation period, mass transfer is mainly driven by the loss of angular momentum."
+ The orbital period will decrease with the donor nass until a minimum period is reached., The orbital period will decrease with the donor mass until a minimum period is reached.
+ Paczvuski&Sienkiewicz(L981) found a minimum period about 80 min without MD. while Podsiadlowskietal.(2002) showed that minimum orbital periods less than 11 min could be reached for binaries with an initial orbital period very close to the bifurcation period if traditional MD is included. but in a time longer (han the age of the universe.," \citet{paczynski81} found a minimum period about $80$ min without MB, while \citet{pod02} showed that minimum orbital periods less than $11$ min could be reached for binaries with an initial orbital period very close to the bifurcation period if traditional MB is included, but in a time longer than the age of the universe."
+" vanderSluvsetal.(2005b) futher imvestigated this ""magnetic capture"" scenario for the formation of UCXNDs."," \citet{sluys05b} further investigated this “magnetic capture"" scenario for the formation of UCXBs."
+ Our calculations show that when the initial orbital period is close to the biftweation period. ulüra-compact svstems (2?«I lh) can indeed form with saturated MD. but also in a (ime longer (han the age of the universe.," Our calculations show that when the initial orbital period is close to the bifurcation period, ultra-compact systems $P<1$ h) can indeed form with saturated MB, but also in a time longer than the age of the universe."
+ For example. for an LAINB with Ma;=L.2M. and 2=0.46 dav in model 4. a final period of £7;=22 min can be reached alter =15 Gyr of mass transfer.," For example, for an LMXB with $M_{2,\rm i}=1.3M_\sun$ and $P_{\rm i}=0.46$ day in model 4, a final period of $P_{\rm f}=22$ min can be reached after $\ga 15$ Gyr of mass transfer."
+ All the works done by previous authors show that a more efficient angular moment loss mechanism is required to produce UCABs within 13.7 αντ in this scenario., All the works done by previous authors show that a more efficient angular momentum loss mechanism is required to produce UCXBs within $13.7$ Gyr in this scenario.
+used the JHK filters and the camera oriented with North up and East to the left.,used the $JHK$ filters and the camera oriented with North up and East to the left.
+" Exposure times were 5 seconds each in an ABBA dither pattern, which was then shifted 10” to the east and repeated (see Figure 3))."," Exposure times were 5 seconds each in an ABBA dither pattern, which was then shifted $\arcsec$ to the east and repeated (see Figure \ref{fig:near-IR}) )."
+ A sky image was created from the 8 images then subtracted from individual frames before relative photometry was performed., A sky image was created from the 8 images then subtracted from individual frames before relative photometry was performed.
+ 2MASS J0850+1057 was targeted in the astrometric programs of ? and ? resulting in two published but discrepant parallax| values., 2MASS J0850+1057 was targeted in the astrometric programs of \citet{2004AJ....127.2948V} and \citet{2002AJ....124.1170D} resulting in two published but discrepant parallax values.
+ The ? work used a near-IR imager with either J or H filters (H in the case of 2MASS J0850--1057) and the ?~~ work used an optical CCD with a wide—I interference filter for astrometric measurements., The \citet{2004AJ....127.2948V} work used a near-IR imager with either $J$ or $H$ filters $H$ in the case of 2MASS J0850+1057) and the \citet{2002AJ....124.1170D} work used an optical CCD with a $wide-I$ interference filter for astrometric measurements.
+" There were seven objects in common between the two programs, five of which had parallax measurements that matched better than 4 mas."," There were seven objects in common between the two programs, five of which had parallax measurements that matched better than 4 mas."
+ 2MASS J0850--1057 was discrepant by 12.9-+5.5 mas (CCD-IR) resulting in a 2.40 difference in parallactic angle and a 9.4σ difference in proper motion position angle., 2MASS J0850+1057 was discrepant by $\pm$ 5.5 mas (CCD-IR) resulting in a $\sigma$ difference in parallactic angle and a $\sigma$ difference in proper motion position angle.
+ ? investigated whether orbital motion between the resolved components of the binary could account for the disagreement but ruled it out due to the short time-span between the two programs compared to the predicted ~51 year orbit., \citet{2004AJ....127.2948V} investigated whether orbital motion between the resolved components of the binary could account for the disagreement but ruled it out due to the short time-span between the two programs compared to the predicted $\sim$ 51 year orbit.
+" We noticed on recent ISPI images that there was a second fainter object resolved in the J band ~ 4.0"" from the L dwarf that could have skewed prior astrometric measurements (see Figure ", We noticed on recent ISPI images that there was a second fainter object resolved in the $J$ band $\sim$ $\arcsec$ from the L dwarf that could have skewed prior astrometric measurements (see Figure \ref{fig:blend}) ).
+"The ? parallax measurement was based on 1.86 4)).years of data, 13 images and a mean epoch of 2001.791."," The \citet{2004AJ....127.2948V} parallax measurement was based on 1.86 years of data, 13 images and a mean epoch of 2001.791."
+" The ? parallax measurement was based on 3.3 years of data, and 30 images."," The \citet{2002AJ....124.1170D} parallax measurement was based on 3.3 years of data, and 30 images."
+" Since no mean epoch is given, we assume observations were conducted between 1999-2002 shortly after the discovery and before the 7 publication."," Since no mean epoch is given, we assume observations were conducted between 1999-2002 shortly after the discovery and before the \citet{2002AJ....124.1170D} publication."
+ Both astrometric programs have comparable resolutions to ISPI., Both astrometric programs have comparable resolutions to ISPI.
+ Tracing the position of 2MASS J0850--1057 back to the ? and ? mean positions from the ISPI position using our updated proper motion vector revealed that the binary L dwarf would have been blended with the contaminant at the time of the cited astrometric programs., Tracing the position of 2MASS J0850+1057 back to the \citet{2004AJ....127.2948V} and \citet{2002AJ....124.1170D} mean positions from the ISPI position using our updated proper motion vector revealed that the binary L dwarf would have been blended with the contaminant at the time of the cited astrometric programs.
+" Over the 1.86-3.3 years of the programs, the significant proper motion of 2MASS J0850+1057 would have changed the centroid shape as its separation with the contaminant increased."," Over the 1.86-3.3 years of the programs, the significant proper motion of 2MASS J0850+1057 would have changed the centroid shape as its separation with the contaminant increased."
+" This potential ""elongation"" of the centroid would have skewed the position of 2MASS J0850--1057 in the frames used to calculate the parallax and proper motion."," This potential ""elongation"" of the centroid would have skewed the position of 2MASS J0850+1057 in the frames used to calculate the parallax and proper motion."
+ A digital sky survey image from 1995 November 17 in J band as well as a Sloan Digital Sky Survey (?)) image from 2005 November 06 in the sloan 7 band also show the point source., A digital sky survey image from 1995 November 17 in $I$ band as well as a Sloan Digital Sky Survey \citealt{2000AJ....120.1579Y}) ) image from 2005 November 06 in the sloan $i$ band also show the point source.
+ The second object appears to be 2 magnitudes fainter in J but shows no appreciable motion and we conclude that it is unrelated to the brown dwarf., The second object appears to be 2 magnitudes fainter in $J$ but shows no appreciable motion and we conclude that it is unrelated to the brown dwarf.
+ We propose that this contaminating source is the most probable cause for the discrepant parallax results., We propose that this contaminating source is the most probable cause for the discrepant parallax results.
+ 2MASS J0850--1057 was imaged as part of the Brown Dwarf Kinematics Project (BDKP) parallax program which targets nearby brown dwarfs., 2MASS J0850+1057 was imaged as part of the Brown Dwarf Kinematics Project (BDKP) parallax program which targets nearby brown dwarfs.
+ We used the Carnegie Astrometric Planet Search software (from here-on ATPa) to extract all point sources in the 6 epochs of ISPI data and solve for the parallax and proper motion (?))., We used the Carnegie Astrometric Planet Search software (from here-on ATPa) to extract all point sources in the 6 epochs of ISPI data and solve for the parallax and proper motion \citealt{2009PASP..121.1218B}) ).
+ The highest quality image was used as the template which all other images were transformed., The highest quality image was used as the template which all other images were transformed.
+" In the 2MASS J0850--1057 field, there were 25 well-behaved (elongated, saturated, spurious sources were removed) reference stars between all epochs."," In the 2MASS J0850+1057 field, there were 25 well-behaved (elongated, saturated, spurious sources were removed) reference stars between all epochs."
+" Using these sources, a linear transformation was applied to each epoch point source catalog to constrain the field rotation, plate scale and match all reference sources to the template."," Using these sources, a linear transformation was applied to each epoch point source catalog to constrain the field rotation, plate scale and match all reference sources to the template."
+ Higher order transformations were tested but demonstrated negligible difference from linear solutions., Higher order transformations were tested but demonstrated negligible difference from linear solutions.
+ The apparent trajectory of each star was fit to a standard astrometric model included in the ATPa software., The apparent trajectory of each star was fit to a standard astrometric model included in the ATPa software.
+ The algorithms follow theastrometric solution prescriptions laid out in the Hipparcos (?)) and Tycho Catalogue (?))tions?.. , The algorithms follow theastrometric solution prescriptions laid out in the Hipparcos \citealt{1997A&A...323L..49P}) ) and Tycho Catalogue \citealt{2000yCat.1259....0H}) ).
+We corrected the data from apparent parallax to absolute parallax following the same procedure described in ?.., We corrected the data from apparent parallax to absolute parallax following the same procedure described in \citet{2004AJ....127.2948V}.
+ A full detailed description of the astrometric pipeline used to solve for the parallax and proper motion of 2MASS J0850--1057 is provided in Faherty et al. (, A full detailed description of the astrometric pipeline used to solve for the parallax and proper motion of 2MASS J0850+1057 is provided in Faherty et al. (
+in ,in prep).
+As noted in Vrba et al (2004) the proper motion and prep).parallax of the system will not be effected by the orbital motion since the images were taken over a ~ 2 year period whereas the orbit is ~ 50 years with components separated by 0.16”., As noted in Vrba et al (2004) the proper motion and parallax of the system will not be effected by the orbital motion since the images were taken over a $\sim$ 2 year period whereas the orbit is $\sim$ 50 years with components separated by $\arcsec$ .
+[or bv increased. absorption which is found in the Miniutti et al. (,for by increased absorption which is found in the Miniutti et al. (
+2009) analvsis.,2009) analysis.
+ Model Fo starts. with Model. Boas the base model but instead. blurring the single reflection spectrum with the convolution model to model any broad residuals present in the spectra., Model F starts with Model B as the base model but instead blurring the single reflection spectrum with the convolution model to model any broad residuals present in the spectra.
+ The beh line complex can still be modelled in. this wav. since the moclel ucles Kika emission and blurring the spectrum. emulates the resulting profile [rom a model., The K line complex can still be modelled in this way since the model includes ${\alpha}$ emission and blurring the spectrum emulates the resulting profile from a model.
+ Since the narrow Felxlxo. core included: within is now relativistically blurred. a narrow Caussian of ixecl width LOecY was added to ensure that the narrow kkeV is still modelled.," Since the narrow ${\alpha}$ core included within is now relativistically blurred, a narrow Gaussian of fixed width eV was added to ensure that the narrow keV is still modelled."
+ This may represent the case. where the kkeV. line is observed from Compton-thin matter. such as the DLIt or NLR.," This may represent the case where the keV line is observed from Compton-thin matter, such as the BLR or NLR."
+ Contrary to Model E. the soft excess (where present) is modelled. using the soft photon Comptonization model.," Contrary to Model E, the soft excess (where present) is modelled using the soft photon Comptonization model."
+ Any donized emission due to and. required. in. Model D. is also found to be required. in this model., Any ionized emission due to and required in Model D is also found to be required in this model.
+ In. general a good fit to all objects is obtained. with Model E clearly. providing a better fit to the data compared το Model. LO for. Fatrall 9. Νας 7469) and SWIET J2127.4|5654.," In general a good fit to all objects is obtained, with Model F clearly providing a better fit to the data compared to Model E for Fairall 9, NGC 7469 and SWIFT J2127.4+5654."
+ The parameters are also consistent within error. bars with the emissivitv index. inclination and spin parameter found. using the line. profile previously. vieleing typically slightly lower emissivity. inclücies.," The parameters are also consistent within error bars with the emissivity index, inclination and spin parameter found using the line profile previously, yielding typically slightly lower emissivity indicies."
+" Note that the best-fitting spin parameter value with Model F for Fairall 9 is à=).40O33ojo (quoted at the confidence level and is unconsrained at the confidence level) in agreement with @=0.44.YO) found in Model D (at the confidence level) and with @=0.5""1 mmfound by Schmoll et al. (", Note that the best-fitting spin parameter value with Model F for Fairall 9 is $a=0.40^{+0.33}_{-0.40}$ (quoted at the confidence level and is unconstrained at the confidence level) in agreement with $a=0.44^{+0.04}_{-0.11}$ found in Model D (at the confidence level) and with $a=0.5^{+0.1}_{-0.3}$ found by Schmoll et al. (
+2009).,2009).
+ Given the independent modelling. of the soft. excess within this model. it is interesting to note that a—0.72.OLLS12 is obtained for NGC 7469 at the confidence level.," Given the independent modelling of the soft excess within this model, it is interesting to note that $a=0.72^{+0.18}_{-0.17}$ is obtained for NGC 7469 at the confidence level."
+ Ehis value of the spin parameter is in agreement with that found in Model D. although providing a slightly worse constraint.," This value of the spin parameter is in agreement with that found in Model D, although providing a slightly worse constraint."
+ Analvsis of the data for these objects obtained with (all objects other than SWIFT. 2127.4|5654. see “Table 2)) vields no additional constraints upon the properties of the Felt region. the reflection component," Analysis of the data for these objects obtained with (all objects other than SWIFT J2127.4+5654, see Table \ref{tab:observations}) ) yields no additional constraints upon the properties of the K region, the reflection component"
+spirals (see Chengalur Ivauckar (2000) and references therein).,spirals (see Chengalur Kanekar (2000) and references therein).
+ Ilieh spin temperatures arise uaturally in cdwarfs because these systems have low metallicities aud pressures and. consequenth. a lavecr fraction of the urna phase of ID as compared to large spiral disks (Cheugalur&IKanuckur 2000).," High spin temperatures arise naturally in dwarfs because these systems have low metallicities and pressures and, consequently, a larger fraction of the warm phase of H as compared to large spiral disks \cite{ck2000}) )."
+ It is possible. of course. that the absorbing galaxies. hough faint iu the optical. are nouctheless exceediuglv eas-rich. and their large II euvolopes cause them to ο preferentially detected in absorption survers.," It is possible, of course, that the absorbing galaxies, though faint in the optical, are nonetheless exceedingly gas-rich, and their large H envelopes cause them to be preferentially detected in absorption surveys."
+ At the very lowest redshifts. the latter bypothesis can be tested hrough deep searches for IT 21 cii emissiou frou the absorbers.," At the very lowest redshifts, the latter hypothesis can be tested through deep searches for H 21 cm emission from the absorbers."
+ Such searches vield direct estimates of the ID nass aud can thus be used to check whether or not the optically faint ealaxies which eive rise to DLA absorption rave anomalously huge TD coutent., Such searches yield direct estimates of the H mass and can thus be used to check whether or not the optically faint galaxies which give rise to DLA absorption have anomalously large H content.
+ We describe. iu lis letter. a deep search for 21 cimi cuuissiou/absorption ron a candidate damped absorber at +~0.101 towards he quasar (Petitjeanetal. 1996)).," We describe, in this letter, a deep search for 21 cm emission/absorption from a candidate damped absorber at $z \sim 0.101$ towards the quasar \cite{petitjean96}) )."
+ The observations were carricd out with the Australia Telescope Compact Arrav (ATCA)., The observations were carried out with the Australia Telescope Compact Array (ATCA).
+ No cussion was detected. resulting in strong constraints on the IT mass of the absorber.," No emission was detected, resulting in strong constraints on the H mass of the absorber."
+ PISS 133 was observed using the 1.5À coufiguration of the ATCA ou a nuuber of occasions in December 1999 and January 2000., PKS $-$ 433 was observed using the 1.5A configuration of the ATCA on a number of occasions in December 1999 and January 2000.
+ The total ou-source mteeration time was ~65 hours., The total on-source integration time was $\sim65$ hours.
+ A bandwidth of δ MIIz was used or the observations. divided mto 1021 channels. aud centred at 1290 MIIz.," A bandwidth of 8 MHz was used for the observations, divided into 1024 channels, and centred at 1290 MHz."
+ This vielded avelocity resolution of ~ 1.8 lan + and a total velocity coverage of ~1800 kin sf., This yielded avelocity resolution of $\sim$ 1.8 km $^{-1}$ and a total velocity coverage of $\sim 1800$ km $^{-1}$.
+ Ouly NN iud YY polarizations were measured., Only XX and YY polarizations were measured.
+ The stroug SOULCC PRS 136 was used for :uplitude/plase and bandpass calibration: this was observed every forty minutes m each observing session., The strong source PKS $-$ 436 was used for amplitude/phase and bandpass calibration; this was observed every forty minutes in each observing session.
+ 638 was used as the primary flux calibrator. and observed at least once ming cach observiug run.," $-$ 638 was used as the primary flux calibrator, and observed at least once during each observing run."
+ The data were analyzed using the software packageMIRIAD., The data were analyzed using the software package.
+ Some baselines of the array were found to eive intermittent correlator errors., Some baselines of the array were found to give intermittent correlator errors.
+ Every 30-second averaged spectrmm was lence mspected separately for cach baseline to detect and reject auv correlator errors as well as to reject any obvious spectral interference., Every 30-second averaged spectrum was hence inspected — separately for each baseline — to detect and reject any correlator errors as well as to reject any obvious spectral interference.
+ Next. the routine UVLIN was used to subtract out continua cnussion from discrete sources in the field.," Next, the routine UVLIN was used to subtract out continuum emission from discrete sources in the field."
+ A threc-dimensional data cube was then made from the residual (i.c. continuum subtracted) visibilitics: spectra or further analyses were constructed usine this cube., A three-dimensional data cube was then made from the residual (i.e. continuum subtracted) visibilities; spectra for further analyses were constructed using this cube.
+ The data from the differcut observing sessions were shited to the helioceutzie frame before being averaged to produce the Spectruni., The data from the different observing sessions were shifted to the heliocentric frame before being averaged to produce the final spectrum.
+ While PKS 183 is unresolved by ic ATCA svutlesized bea. the field also contains the uuch stronecr compact source PISS 136 (flux ~Lit Jy. our observatious). near the edge of the primary beam. as well as oue more weak conrpact source (flux 7 LOO ταν. our observations).," While PKS $-$ 433 is unresolved by the ATCA synthesized beam, the field also contains the much stronger compact source PKS $-$ 436 (flux $\sim 
+4.14$ Jy, our observations), near the edge of the primary beam, as well as one more weak compact source (flux $\sim$ 400 mJy, our observations)."
+ It was thus vossible that the use of UVLIN for σολΙΙ straction miehlt result in spectral artefacts (Coruwelletal. 19923)., It was thus possible that the use of UVLIN for continuum subtraction might result in spectral artefacts \cite{cornwell}) ).
+ We hence also tied an alternate analvsis procedure to check for spectral errors that might arise due to the confusing sources in the field., We hence also tried an alternate analysis procedure to check for spectral errors that might arise due to the confusing sources in the field.
+ This involved the use of the task UVSUD. to subtract out the individual point sources: UVLIN was then run to remove any residual continuum cussion aud the resulting data set was again mapped to obtain the spectral data cube.," This involved the use of the task UVSUB, to subtract out the individual point sources; UVLIN was then run to remove any residual continuum emission and the resulting data set was again mapped to obtain the spectral data cube."
+ The spectra obtained frou the two procedures were found to be identical within the noise for cach data set., The spectra obtained from the two procedures were found to be identical within the noise for each data set.
+ Au RAIS noise of 1.56 wy was obtained per chamuel at the original velocity resolution of 1.8 kan +: this is close to the theoretical scusitivity of the ATCA. for our observing parameters.," An RMS noise of 1.56 mJy was obtained per channel at the original velocity resolution of 1.8 km $^{-1}$; this is close to the theoretical sensitivity of the ATCA, for our observing parameters."
+ Figure d. shows the spectrum smoothed to a resolution of 9 kins |: weal: absorption call ρα seen Close to 2=0.101., Figure \ref{fig:fig2} shows the spectrum smoothed to a resolution of 9 km $^{-1}$: weak absorption can be seen close to $z = 0.101$.
+ This occurs at a helocentiic frequency of 1290.1LL MITz. corresponding to a redshift of 0.10097+0.000023: this is in good aereemoeut with the redshift 2=0.101 obtained from metal lines (Petitjeanetal. 1996)).," This occurs at a heliocentric frequency of 1290.144 MHz, corresponding to a redshift of $0.10097 \pm 0.00003$; this is in good agreement with the redshift $z = 0.101$ obtained from metal lines \cite{petitjean96}) )."
+ The RAIS noise ou the spectrum is ~ 0.76 inJv while the feature is ~ 2.5 mJv. deep (aud oulv one chanuel wide). ie. a 3.30 result.," The RMS noise on the spectrum is $\sim$ 0.76 mJy while the feature is $\sim$ 2.5 mJy deep (and only one channel wide), i.e. a $\sigma$ result."
+ The absorption was seen in both he XX and YY polarizations separately. although. of course. at even lower significance levels.," The absorption was seen in both the XX and YY polarizations separately, although, of course, at even lower significance levels."
+ Figure 2 shows a zoomed-in version of the spectrum at the original resolution of 1.8 km |: the feature of fieure can be scen to be a few channels wide here (although within the noise)., Figure \ref{fig:fig1} shows a zoomed-in version of the spectrum at the original resolution of 1.8 km $^{-1}$ ; the feature of figure \ref{fig:fig2} can be seen to be a few channels wide here (although within the noise).
+ We note that the narrowness of the, We note that the narrowness of the
+(Québec).,.
+.. JL acknowledges support from the National Science Foundation (through grant. AST-0307321 [or study of white dwarfs., JL acknowledges support from the National Science Foundation through grant AST-0307321 for study of white dwarfs.
+ UIXIRT. the United Kingdom Infrared. Telescope. is operated by the Joint Astronomy. Centre on behalf of the U.IX. Particle Physies and Astronomy Research Council.," UKIRT, the United Kingdom Infrared Telescope, is operated by the Joint Astronomy Centre on behalf of the U.K. Particle Physics and Astronomy Research Council."
+propagating up with an amplitude of hits the shock. it creates an advected wave with an amplitude 5/7 as well as an acoustic wave propagating down with an amplitude f°®.,"propagating up with an amplitude $\delta f^{-}$ hits the shock, it creates an advected wave with an amplitude $\delta f^{\rm adv}$ as well as an acoustic wave propagating down with an amplitude $\delta f^{+}$."
+" The ratio of the amplitude of the waves created to iat of the incoming wave defines the coupling constants: Qa[4—of""fof and Ry=of.fof.", The ratio of the amplitude of the waves created to that of the incoming wave defines the coupling constants: $Q_\sh \equiv \delta f^{\rm adv}/\delta f^-$ and $R_\sh \equiv \delta f^{+}/\delta f^-$.
+ ος can be 'omputed. analytically using the boundary conditions at je shock and by separating the perturbations into the Ίος tvpes of waves owing to the WIXD method (??)..," These can be computed analytically using the boundary conditions at the shock and by separating the perturbations into the three types of waves owing to the WKB method \citep{foglizzo05,foglizzo09}."
+ When an entropy-vorticity wave is acdvected towards ιο PNS. it creates an acoustic wave propagating up with in ellicieney. characterised by the ratio of the amplitudes (measured at the shock): Qc=8f.fof.," When an entropy-vorticity wave is advected towards the PNS, it creates an acoustic wave propagating up with an efficiency characterised by the ratio of the amplitudes (measured at the shock): $Q_\nabla \equiv \delta f^{-}/\delta f^{\rm adv}$."
+ Similarly: ie rellection in the gradients of the acoustic wave oopagating down is characterised bv: Ac—affof, Similarly the reflection in the gradients of the acoustic wave propagating down is characterised by: $R_\nabla \equiv \delta f^{-}/\delta f^{+}$.
+ These two constants can be computed. numerically ov integrating the cillerential svstem describing the »erturbations. imposing mocdified boundary conditions at he shock (9f=0 to obtain Ως. af=0 to obtain Hx).," These two constants can be computed numerically by integrating the differential system describing the perturbations, imposing modified boundary conditions at the shock $\delta f^+=0$ to obtain $Q_\nabla$, $\delta f^{\rm adv} = 0$ to obtain $R_\nabla$ )."
+ Finally. using these four coupling constants. one can celine the elobal constants that describe the ellicicney of the two eveles (the amplitude ratio between the beginning and the end of the eveles): Q=Q.Qx and RoBate.," Finally, using these four coupling constants, one can define the global constants that describe the efficiency of the two cycles (the amplitude ratio between the beginning and the end of the cycles): $Q \equiv Q_\sh Q_\nabla$ and $R \equiv R_\sh R_\nabla$."
+ The evele constants are complex numbers that depend on the complex frequeney c., The cycle constants are complex numbers that depend on the complex frequency $\omega$ .
+ Since the two eveles take place simultaneously. an cigenmode must satisly the following relation (?):: If these constants are calculated with a real frequeney then the instability eriterion for a single evcle can be written: [CQ]2 Lor |J?]71 on (dependingthe evele considered). which is equivalent to the criterion stated in Section 2.2..," Since the two cycles take place simultaneously, an eigenmode must satisfy the following relation \citep{foglizzo00}: If these constants are calculated with a real frequency then the instability criterion for a single cycle can be written: $|Q|>1$ or $|R|>1$ (depending on the cycle considered), which is equivalent to the criterion stated in Section \ref{sec:review_mechanisms}."
+ The growth rate can then be approximated by: where r is the duration of the evele considered (here the advyective-acoustic one)., The growth rate can then be approximated by: where $\tau$ is the duration of the cycle considered (here the advective-acoustic one).
+ When a second. evele is present. it can interfere constructively or destructively with the first evele depending on their relative phase.," When a second cycle is present, it can interfere constructively or destructively with the first cycle depending on their relative phase."
+ ‘This causes oscillations of the growth rate. which are indeed. observed in the cigenspectrum of SASL (Figure 7 of ?)).," This causes oscillations of the growth rate, which are indeed observed in the eigenspectrum of SASI (Figure 7 of \citet{foglizzo07}) )."
+ These oscillations contain information about the ellicicney and. the timescale of the two ceveles at work., These oscillations contain information about the efficiency and the timescale of the two cycles at work.
+ 7. could thus extract directly from the eigenspecteum. the values of the constants Q and /? (note that this is independent. from the WIND analysis. and only assumes the presence of two eveles).," \citet{foglizzo07} could thus extract directly from the eigenspectrum the values of the constants $Q$ and $R$ (note that this is independent from the WKB analysis, and only assumes the presence of two cycles)."
+ Phe value. of the evele constants obtained from the WIXD method and from the iinalvsis of the eigenspectrum are in very good agreement heir Figures δ and 9)., The value of the cycle constants obtained from the WKB method and from the analysis of the eigenspectrum are in very good agreement (their Figures 8 and 9).
+ They show that the advective-acoustic can be unstable with a evele constant reaching Q~6. while the purely. acoustic evele is always stable with tvpically /?0.5.," They show that the advective-acoustic can be unstable with a cycle constant reaching $Q \sim 6$, while the purely acoustic cycle is always stable with typically $R\sim0.5$."
+ The WKB analysis of ? sununarisecd in the last subsection has proven that the adyective-acoustic cvcle is unstable anc can explain SASL. while the purely acoustic evele is stable.," The WKB analysis of \citet{foglizzo07} summarised in the last subsection has proven that the advective-acoustic cycle is unstable and can explain SASI, while the purely acoustic cycle is stable."
+ However. this result. is. valid only for high frequency. modes. satisfying Equation 4..," However this result is valid only for high frequency modes, satisfying Equation \ref{eq:WKB_criterion}."
+ The first few harmonics do not satisfy. this inequality as αμοL lor the fundamental mode., The first few harmonics do not satisfy this inequality as $\omega r_\sh/c_\sh \sim 1$ for the fundamental mode.
+ One may consider that the WIND method can be safely applied for the tenth harmonics and higher (satisfving crac10)., One may consider that the WKB method can be safely applied for the tenth harmonics and higher (satisfying $\omega r_\sh/c_\sh > 10$ ).
+ Unfortunately these high frequency modes. are unstable only for lage shock radi ονc510 (depending on the cooling Function chosen)., Unfortunately these high frequency modes are unstable only for large shock radii $r_\sh/r_*>5-10$ (depending on the cooling function chosen).
+ Γιαν strictly speaking in the more realistic range of shock radius 2€nur<5. the WKB analysis cannot conclude on the instability mechanism.," Thus strictly speaking in the more realistic range of shock radius $2< r_\sh/r_*<5$, the WKB analysis cannot conclude on the instability mechanism."
+ ? have however argued by a continuity argument that it is not necessary to invoke a different instability mechanism for the low frequency modes., \citet{foglizzo07} have however argued by a continuity argument that it is not necessary to invoke a different instability mechanism for the low frequency modes.
+ Indeed. the cigenspectrum vary smoothly. both with the shock radius ancl mode frequency. ancl the cigenfunetion of the low frequeney modes resembles tha of the high frequency modes.," Indeed, the eigenspectrum vary smoothly both with the shock radius and mode frequency, and the eigenfunction of the low frequency modes resembles that of the high frequency modes."
+ It may. thus seem natural to assume that all SASL moces are due to the advective-acoustic evele., It may thus seem natural to assume that all SASI modes are due to the advective-acoustic cycle.
+ Another argument in this direction is tha the WIXD approximation usually gives quite good results even when the small parameter (here eg/(row)) is of order unity., Another argument in this direction is that the WKB approximation usually gives quite good results even when the small parameter (here $c_\sh/(r_\sh\omega)$ ) is of order unity.
+ This is however not a formal proof. anc it cannot be excluded that the low frequency modes al moderate shock radius originate [rom a cdilleren instability mechanism.," This is however not a formal proof, and it cannot be excluded that the low frequency modes at moderate shock radius originate from a different instability mechanism."
+ ? has studied analytically the stability. of a spherical accretion shock hy establishing an approximate dispersion relation. using a method inspired. from ?..," \citet{laming07} has studied analytically the stability of a spherical accretion shock by establishing an approximate dispersion relation, using a method inspired from \citet{vishniac89}."
+ This equation contains terms due to advection. which ? proposes to remove in order to assess the importance of the advective-acoustic evcle in the instability mechanism.," This equation contains terms due to advection, which \citet{laming07} proposes to remove in order to assess the importance of the advective-acoustic cycle in the instability mechanism."
+ Vhe conclusion he initially reached by this mean depended on the radius of the shock., The conclusion he initially reached by this mean depended on the radius of the shock.
+ At large shock radii. the omission of the acvective term. stabilisecl the How. thus favouring the acvective-acoustic evele.," At large shock radii, the omission of the advective term stabilised the flow, thus favouring the advective-acoustic cycle."
+ On the contrary. at small shock radii advection allected very Little the growth rate. flavouring a purely acoustic interpretation of the instabilitv.," On the contrary, at small shock radii advection affected very little the growth rate, favouring a purely acoustic interpretation of the instability."
+ Llowever. this analysis has been criticised by 7. because the dispersion relation was established with the use of dubious approximations.," However, this analysis has been criticised by \citet{yamasaki08} because the dispersion relation was established with the use of dubious approximations."
+ Furthermore an erratum published later on (72). corrects a [ew mistakes and changes the conclusion: the advective- evele is then favoured for all shock ασ., Furthermore an erratum published later on \citep{laming08} corrects a few mistakes and changes the conclusion: the advective-acoustic cycle is then favoured for all shock radii.
+ A number of authors have tried (ο distinguish the two mechanisms by using the oscillation [requenevy observed in the simulations., A number of authors have tried to distinguish the two mechanisms by using the oscillation frequency observed in the simulations.
+ These studies have lec to cliverging conclusions: ? conclucled that the adveetive- evcle is responsible for the instability. while ? concluded in the favour of the purely acoustic mechanism.," These studies have led to diverging conclusions: \citet{ohnishi06} concluded that the advective-acoustic cycle is responsible for the instability, while \citet{blondin06} concluded in the favour of the purely acoustic mechanism."
+ This confusion probably comes fromthe fact that both of the mechanisms can explain the oscillation frequency. if one makes the right assumptions., This confusion probably comes fromthe fact that both of the mechanisms can explain the oscillation frequency if one makes the right assumptions.
+ Indeed. ? exelucled the," Indeed, \citet{ohnishi06} excluded the"
+"Recent works on coustrainiug the cosmological parameters using the CMB atid the high redshift supernovae seein to converge to a “staucdard cosmological model favouring a flat universe will 0,,~0.3 and 03~ Q.T: White '..",Recent works on constraining the cosmological parameters using the CMB and the high redshift supernovae seem to converge to a “standard cosmological model” favouring a flat universe with $\Omega_m\sim 0.3$ and $\Omega_\lambda\sim 0.7$ : White \cite{White}.
+" However these results are still uncertain axl cepeu« OL sole physical assuimptioIs. SO he flat £2,,=1 1uodel is still possible (Le Dour L73)."," However these results are still uncertain and depend on some physical assumptions, so the flat $\Omega_m=1$ model is still possible (Le Dour \cite{LeDour}) )."
+" It is therefore important to explo'e other independent techniques to coustrain these cosmological pa""alueters.", It is therefore important to explore other independent techniques to constrain these cosmological parameters.
+ In cluster gravitationalal leusiug. the existence of multiple images — wih known redshifts — given by the saue sourceUx allows to calibrate in an absolute way the total cluster mass deduced from the eus model.," In cluster gravitational lensing, the existence of multiple images – with known redshifts – given by the same source allows to calibrate in an absolute way the total cluster mass deduced from the lens model."
+ The great improvernuent tn tle mass mocleliig of cluster-lenses that iucludes the cluster galaxies⋅ halos (Ixneib∣⊳⋅ ⋅↽≻aL7.. Natarajan-. -. Iueib ) leads to the hope that clusters cau aso be used to coustrain the geometry of tlie Universe. twoughthe ratio of," The great improvement in the mass modeling of cluster-lenses that includes the cluster galaxies halos (Kneib \cite{Kneib96}, Natarajan Kneib \cite{Natarajan}) ) leads to the hope that clusters can also be used to constrain the geometry of the Universe, throughthe ratio of"
+niereers. rather than galactic protospheroids.,"mergers, rather than galactic protospheroids."
+ Since the nuniber density of the brightest of such objects may be more a funcion of the probalilitv and duration of such starbursts rather than the nature of he underline cosmological model. it may be more useful to use the star formation or metal injection rates [15] indicated by the toal observed rest-frame uItraviolet light to constrain models 16]..," Since the number density of the brightest of such objects may be more a function of the probability and duration of such starbursts rather than the nature of the underlying cosmological model, it may be more useful to use the star formation or metal injection rates \cite{Madau} indicated by the total observed rest-frame ultraviolet light to constrain models \cite{SomPri}."
+ The available data ou the historv of sar formation [77J7S.το). suggests hat most of the stars aud most of the metals observed formed relatively receut» after about τοςIshif ~|: and that the toal star formation rate at ~Qi 19 rhaps a factor of3 lower than at  WI1hn ve another factor of ~EMi aloft to 1 (although the rates atf olAmo3 CDu 0 higher if most of f star formaion is iu objects too faint to see}.," The available data on the history of star formation \cite{Gallego,Lilly,Madau} suggests that most of the stars and most of the metals observed formed relatively recently, after about redshift $z\sim 1$; and that the total star formation rate at $z\sim 3$ is perhaps a factor of 3 lower than at $z \sim 3$ with yet another factor of $\sim 3$ falloff to $z \sim 4$ (although the rates at $z\gsim 3$ could be higher if most of the star formation is in objects too faint to see)."
+ This is in accord with indicatio1s ron dauped Ly: ula SVenis [79]. and expecta10115 or Q=] inodels such as CIIDM. lnt not wih the expectations for low-Qy nodes whicli ave less crowhi of fluctiations at recent epochs. and therefore mast form strucure earlier.," This is in accord with indications from damped Lyman $\alpha$ systems \cite{FallCP} and expectations for $\Omega=1$ models such as CHDM, but not with the expectations for $\Omega_0$ models which have less growth of fluctuations at recent epochs, and therefore must form structure earlier."
+ Dut lis waist o investigated 1sine more detailed modeling. incliding gas cooling aud feeccback from stars auid supernovae [76].. before strong couclusious cau be dran.," But this must be investigated using more detailed modelling, including gas cooling and feedback from stars and supernovae \cite{SomPri}, before strong conclusions can be drawn."
+ There is another sort oconstraint from observed muubers of ligh-redshitt protogaaxies that would appear to disfavorACDM., There is another sort of constraint from observed numbers of high-redshift protogalaxies that would appear to disfavor.
+. The upper luit ou the nunber of:& Lobjects in the Hubble Deep Field (whic1 presuinably correspond to stallcr-luass Galaxies than most of the Steidel objects) is far lower than tie expectaions in m models. especially with a positive cosmological coustaut. because of the large volume at high redshift iu such cosinologies [80]..," The upper limit on the number of $z\gsim 4$ objects in the Hubble Deep Field (which presumably correspond to smaller-mass galaxies than most of the Steidel objects) is far lower than the expectations in $\Omega_0$ models, especially with a positive cosmological constant, because of the large volume at high redshift in such cosmologies \cite{Yahil}."
+ Thus evidence from high-redslüft objects cuts both wavs. and it is too early to tell whether hieh- or low-Qy mocels will ultimately be favored.," Thus evidence from high-redshift objects cuts both ways, and it is too early to tell whether high- or $\Omega_0$ models will ultimately be favored."
+ There are three basic reasous why a iuixture of cold plus hot dark matter works setter than pure CDM without any hot particles: the power spectrum shape ΙΛ) Ίνα etter fit to observations. there are indications from observations OY a niore weakly clustering componeit of dark matter. and a hot componcut nav help avoid the too-deuse ceutra dark matter density 1ji pure CDAL dark uatter haos.," There are three basic reasons why a mixture of cold plus hot dark matter works better than pure CDM without any hot particles: the power spectrum shape $P(k)$ is a better fit to observations, there are indications from observations for a more weakly clustering component of dark matter, and a hot component may help avoid the too-dense central dark matter density in pure CDM dark matter halos."
+ T will discuss cac iin turn., I will discuss each in turn.
+" As I explained im discussing WDM vs, CIIDM above. he pure CDA spectrum Ph) does not fall fast οποιο] on he large-A side of its peak il oreler to fit iudicaions from galaxy and cluster correlations aud ower spectra."," As I explained in discussing WDM vs. CHDM above, the pure CDM spectrum $P(k)$ does not fall fast enough on the $k$ side of its peak in order to fit indications from galaxy and cluster correlations and power spectra."
+ The discussion here of “Excess Power is à wav of quautifving lis., The discussion there of “Excess Power” is a way of quantifying this.
+ This is aso related to tli' overproduction of clusters i pure CDAL, This is also related to the overproduction of clusters in pure CDM.
+" The obvious wav to prevent ο= 1xCDAL normalized to CODE fron overproduciug clusters is to fit it a lot (the xecise amount depending ou how muuch of the CODE ductualous are attribued to eravity waves,which cau be increasingly nuportaut astje tilt is iucreased)."," The obvious way to prevent $\Omega=1$ sCDM normalized to COBE from overproducing clusters is to tilt it a lot (the precise amount depending on how much of the COBE fluctuations are attributed to gravity waves,which can be increasingly important as the tilt is increased)."
+ But a coustraiut on CDM-tvpo models that is, But a constraint on CDM-type models that is
+(solid. dashed. dotted and dashed-dotted Hues respectively for increasingly smaller blocks).,"(solid, dashed, dotted and dashed-dotted lines respectively for increasingly smaller blocks)."
+ The difference due to the block size used for resampling appear to be small., The difference due to the block size used for resampling appear to be small.
+ As mentioned. (his was an advantage of the marked point bootstrap.," As mentioned, this was an advantage of the marked point bootstrap."
+ Loh&Stein(2004) found. greater variation of performance with block size for subsample and block bootstrap., \citet{loh02a} found greater variation of performance with block size for subsampling and block bootstrap.
+ Compared with the Poisson empirical coverage. we find that at low densities and smaller observation region sizes (plots towards the upper left of Figure 1)). the bootstrap method does poorly.," Compared with the Poisson empirical coverage, we find that at low densities and smaller observation region sizes (plots towards the upper left of Figure \ref{fig:HamPoiCoverage}) ), the bootstrap method does poorly."
+ However. the empirical coverage of (he bootstrap confidence intervals quickly increases towards with increasing density (down the columns in Figure 1)) and/or observation region size (across the rows in Figure 1)). Le. with larger sample sizes.," However, the empirical coverage of the bootstrap confidence intervals quickly increases towards with increasing density (down the columns in Figure \ref{fig:HamPoiCoverage}) ) and/or observation region size (across the rows in Figure \ref{fig:HamPoiCoverage}) ), i.e. with larger sample sizes."
+continuous line in Figure 2.,continuous line in Figure 2.
+ The dashed lines represent the lo deviations around the mean., The dashed lines represent the $1\sigma$ deviations around the mean.
+ To improve the statistics we consider pairs in the range aaresec. where our estimator is more sensitive in identifving candidate interacting svstenis (sce section. ?27)).," To improve the statistics we consider pairs in the range arcsec, where our estimator is more sensitive in identifying candidate interacting systems (see section \ref{sec_6}) )."
+" From the mock catalogues we count a total of 25.64 5.0. pairs around ""normal optically selected galaxies for angular separations 4«00612 aarcsec and for DPx0.05.", From the mock catalogues we count a total of $\pm$ 5.0 pairs around `normal' optically selected galaxies for angular separations $4<\delta\theta< 12$ arcsec and for $P\le0.05$.
+ This compares to 43.0d6.6 (Poisson statistics) pairs around. radio sources., This compares to $43.0\pm6.6$ (Poisson statistics) pairs around radio sources.
+ Therefore. the nuniber of pairs around faint radio sources above the random and w(A) expectation is (43.0+6.6)(25.645.0)=17448.3 (2.10 confidence level).," Therefore, the number of pairs around faint radio sources above the random and $w(\theta)$ expectation is $(43.0\pm6.6)-(25.6\pm5.0)=17.4\pm8.3$ $\sigma$ confidence level)."
+ A similar result. with a slightly larger significance (2.30). is obtained for P?0.10.," A similar result, with a slightly larger significance $\sigma$ ), is obtained for $P\le0.10$."
+ The environment of [aint radio sources has been explored in previous studies using visual inspection., The environment of faint radio sources has been explored in previous studies using visual inspection.
+ Kron et al. (, Kron et al. (
+1985) considered a radio selected sample with Sy mm. and available four-band photometry (C. J. £F. IN-,"1985) considered a radio selected sample with $S_{1.4}>0.6$ mJy and available four-band photometry $U$ , $J$, $F$, $N$ -bands)."
+ For the sub-sample with 17.0<2«21.0 mmag (hore we assume the transformation {ἱ--0.14 between F' and -band magnitudes: Metcalfe et al., For the sub-sample with $17.0<R<21.0$ mag (here we assume the transformation $R=F-0.14$ between $F$ and $R$ -band magnitudes; Metcalfe et al.
+ 1991) they found hat about 19% of the sources (15 out of 95) lie in. pairs or groups., 1991) they found that about $19\%$ of the sources (18 out of 95) lie in pairs or groups.
+ More recently. Cirappioni et al. (," More recently, Gruppioni et al. ("
+1999) studied a radio sample with Sy2 O.2mmJv and multi-waveleneth yhotometric data.,1999) studied a radio sample with $S_{1.4}>0.2$ mJy and multi-wavelength photometric data.
+ They conclude that about 2 of the sources with Spo)> O.d3mmJy and 17.0</21.0 mmasg (6 out of 22 galaxies in that study) are found in pairs or small groups., They conclude that about $27\%$ of the sources with $S_{1.4}>0.4$ mJy and $17.0<R<21.0$ mag (6 out of 22 galaxies in that study) are found in pairs or small groups.
+ The analysis described in the previous section identified. 43 candidate interacting pairs out of a total of 206 sources. corresponding to a pairing fraction. of 21%. in reasonable agreement with the results from the above mentioned studies.," The analysis described in the previous section identified 43 candidate interacting pairs out of a total of 206 sources, corresponding to a pairing fraction of $21\%$ , in reasonable agreement with the results from the above mentioned studies."
+" We also find that although the pairing fraction of the sub-mJvy. sources in the present sample (5,4cO4 mmy: 17.0«xHzx21.0 mmag: 662 taaresec) is higher than that of optically selected. galaxies (at the same magnitude limit). the significance of the excess is low (s 2a)."," We also find that although the pairing fraction of the sub-mJy sources in the present sample $S_{1.4}\ge0.4$ mJy; $17.0\le R \le 21.0$ mag; $\delta\theta\ge4$ arcsec) is higher than that of optically selected galaxies (at the same magnitude limit), the significance of the excess is low $\approx2\sigma$ )."
+ This implies that the pairing fraction of the sub-m.Jv population is higher than that of local galaxies., This implies that the pairing fraction of the sub-mJy population is higher than that of local galaxies.
+ This is because the frequency of interactions in optically selected: samples has been shown to increase with redshift (Burkey ct al., This is because the frequency of interactions in optically selected samples has been shown to increase with redshift (Burkey et al.
+ 1994)., 1994).
+ Therefore. the optically selected: galaxy sample studied here. with a magnitude limit /?=21.0 mmaeg. corresponding to a mean redshift zz0.25 (Pozzetti. Bruzual Zamorani 1996). already. comprises a higher fraction of interacting systems than local galaxy samples.," Therefore, the optically selected galaxy sample studied here, with a magnitude limit $R=21.0$ mag, corresponding to a mean redshift $z\approx0.25$ (Pozzetti, Bruzual Zamorani 1996), already comprises a higher fraction of interacting systems than local galaxy samples."
+ Moreover. the Lgsimilarity. of the sub-mJy ancl field galaxy close pair distributions in Figure 2 suggests that radio selection at a given optical magnitude limit does not guarantee a significantly higher fraction of interacting systems compared to optical selection (at the same limiting maenitucle)," Moreover, the similarity of the sub-mJy and `field' galaxy close pair distributions in Figure \ref{fig_6.3} suggests that radio selection at a given optical magnitude limit does not guarantee a significantly higher fraction of interacting systems compared to optical selection (at the same limiting magnitude)."
+ llowever. one shoulki be cautious about τὶς interpretation.," However, one should be cautious about this interpretation."
+ Firstlv. the statistical approach used. here calculates the probability of a galaxy having a physically associated companion with projected separation 96 and magnitude A? by estimating the expected number of galaxies xighter than # within radius 98.," Firstly, the statistical approach used here calculates the probability of a galaxy having a physically associated companion with projected separation $\delta\theta$ and magnitude $R$ by estimating the expected number of galaxies brighter than $R$ within radius $\delta\theta$."
+ Therefore. for à given magnitude. 2. the larger the separation. 9006. the higher the xobabilitv. P. that the pair is à spurious alignment of the sky plane.," Therefore, for a given magnitude, $R$, the larger the separation, $\delta\theta$, the higher the probability, $P$, that the pair is a spurious alignment of the sky plane."
+ For example. for /?zz20.0 mimag only pairs with 86]2aarcsec are identified as potentially. interacting or P«0.05.," For example, for $R\approx20.0$ mag only pairs with $\delta\theta\lesssim12$ arcsec are identified as potentially interacting for $P<0.05$."
+ Consequently. the estimator in. equation 2) ) is sensitive to relatively close pairs (6812 aaresee) and is likely to miss physically associated galaxies wit-— arecr projected separations.," Consequently, the estimator in equation \ref{eq_6.2}) ) is sensitive to relatively close pairs $\delta\theta\lesssim12$ arcsec) and is likely to miss physically associated galaxies with larger projected separations."
+ Additionally. due to resolution ellects we can only identifv candidate interacting svstenis with 38laarcsec.," Additionally, due to resolution effects we can only identify candidate interacting systems with $\delta\theta\gtrsim4$ arcsec."
+" This angularὃν separation corresponds o a linear size of 20kpc (44,=50kms‘Alpety αἱ >=0.25. the median recshift of the sub-sample with So.c Odmmlw. 170<Ax210mmag and availab spectroscopic information (79 sources out of a total of 206)."," This angular separation corresponds to a linear size of $20\,\mathrm{kpc}$ $H_{0}=50\mathrm{\,km\,s^{-1}\,Mpc^{-1}}$ ) at $z=0.25$, the median redshift of the sub-sample with $S_{1.4}\ge0.4$ mJy, $17.0\le R\le21.0$ mag and available spectroscopic information (79 sources out of a total of 206)."
+ These very close pairs are also expected {ο have high radio luminosity (normalised to optical luminosity: ltead Ponman 1998)., These very close pairs are also expected to have high radio luminosity (normalised to optical luminosity; Read Ponman 1998).
+ Future high-resolution photometric observations have the potential to reveal the presence of very close pairs within the sub-m.Jv population. that remain unresolved in the present study.," Future high-resolution photometric observations have the potential to reveal the presence of very close pairs within the sub-mJy population, that remain unresolved in the present study."
+ Aloreover. in our analysis we only consider the opticalA xightest (/?—21.0 mmag) radio sources.," Moreover, in our analysis we only consider the optically brightest $R\le21.0$ mag) radio sources."
+ These are likely o lie. on average. at. relatively lower redshifts. compared o optically fainter objects (CH?21.0mmae).," These are likely to lie, on average, at relatively lower redshifts compared to optically fainter objects $R>21.0$ mag)."
+ The pairing raction of the optically fainter sub-niJy. sources reniains to » explored., The pairing fraction of the optically fainter sub-mJy sources remains to be explored.
+ A total of 27 out of 43 (zz60%) faint radio sources or which our analysis indicates possible association with ‘field’ galaxies (I? 0.05). have available spectroscopic information.," A total of 27 out of 43 $\approx60\%$ ) faint radio sources for which our analysis indicates possible association with `field' galaxies $P\le0.05$ ), have available spectroscopic information."
+ The sample consists of absorption-line svstenis C). Sevfert 10r 2 type objects (15%). unclassified sources (3054) and star-Lorming galaxies (19%).," The sample consists of absorption-line systems $37\%$ ), Seyfert 1 or 2 type objects $15\%$ ), unclassified sources $30\%$ ) and star-forming galaxies $19\%$ )."
+ For Px0.10. the relative fraction of dillerent tvpes of sources is similar.," For $P\le0.10$, the relative fraction of different types of sources is similar."
+ The radio power distributions of these sources are compared in Figure 3.— with those of the spectroscopic sample of optically identified: radio galaxies with Sy)2 0.4mmJy and 17.0<Ac 2L0mmag., The radio power distributions of these sources are compared in Figure \ref{fig_6.4} with those of the spectroscopic sample of optically identified radio galaxies with $S_{1.4}\ge0.4$ mJy and $17.0\le R\le 21.0$ mag.
+ The histogram [or the La luminosity distribution of the two samples is shown in Figure 4.., The histogram for the $\alpha$ luminosity distribution of the two samples is shown in Figure \ref{fig_6.5}.
+ There is no significant cillerence between the two populations in these Figures., There is no significant difference between the two populations in these Figures.
+ Hummel (1981). also. founc no dillerence in the radio power distributions of isolate ealaxies and galaxies in pairs., Hummel (1981) also found no difference in the radio power distributions of isolated galaxies and galaxies in pairs.
+ Moreover. interacting svstenis. although being on average more luminous at infrarec or llo. wavelengths compared to Ποια galaxies. exhibi significant scatter around. their mean emission. properties (Ixcel et al.," Moreover, interacting systems, although being on average more luminous at infrared or $\alpha$ wavelengths compared to `field' galaxies, exhibit significant scatter around their mean emission properties (Keel et al."
+ 1985: Ixennicuttet al., 1985; Kennicuttet al.
+ LOST: Bushouse et al., 1987; Bushouse et al.
+ 1985), 1988).
+ This isattributed to (0) the relative velocities of the interacting galaxies (ii) the properties of the ealactic disks. especially in relative weak encounters and (ii) the time scale of the induced star-formation relative to that of the gravitational encounter (Ixennicutt et al.," This isattributed to (i) the relative velocities of the interacting galaxies (ii) the properties of the galactic disks, especially in relative weak encounters and (iii) the time scale of the induced star-formation relative to that of the gravitational encounter (Kennicutt et al."
+ LOST)., 1987).
+ Hence.," Hence,"
+in broad soft. X-ray bands.,in broad soft X-ray bands.
+ In directions where the Galactie disk blocks all of the η keV emission from bevond the Local Bubble. the observedROSAT R2 to R1 band ratio is 1.06. implving a temperature of 10°! IK (INXuntz&Snowden2000).," In directions where the Galactic disk blocks all of the $\frac{1}{4}$ keV emission from beyond the Local Bubble, the observed R2 to R1 band ratio is $\sim 1.06$ , implying a temperature of $10^{6.1}$ K \citep{kuntz_snowden}."
+. This temperature. is consistent with values found by shadowing observations (described below) at higher Galactic latitudes (INuntz&Snowden2000:Snowclenetal.2000) and with Wisconsins (UXT) and All-Skv Survey's B/Be and C/B band ratios (Blochetal.1936:Juda 1991).," This temperature is consistent with values found by shadowing observations (described below) at higher Galactic latitudes \citep{kuntz_snowden,snowden_etal_00}
+ and with Wisconsin's ) and All-Sky Survey's B/Be and C/B band ratios \citep{bloch_etal,juda_etal}."
+. Similar temperatures were also derived [rom the Spectrometers soft N-ray. spectra. although these hieh resolution spectra were difficult to fit. with simple models (Sandersetal.2001).," Similar temperatures were also derived from the 's soft X-ray spectra, although these high resolution spectra were difficult to fit with simple models \citep{sanders_etal}."
+. Desirous of a better understanding of our environs as well as a model of the bubble population (useful lor calculating the hot gas filling factor). both observers ancl (theorists have long sought explanations of the LD's origin and life history.," Desirous of a better understanding of our environs as well as a model of the bubble population (useful for calculating the hot gas filling factor), both observers and theorists have long sought explanations of the LB's origin and life history."
+ Many. researchers have pointed out that the enerev embodied in the Local Bubbles X-ray emitting plasma is so ereal as lo require a supernova or a series of supernovae. possibly aided by stellar winds.," Many researchers have pointed out that the energy embodied in the Local Bubble's X-ray emitting plasma is so great as to require a supernova or a series of supernovae, possibly aided by stellar winds."
+ The energy may have been deposited within an existing gap in the II I disk or may have pushecl aside the disk material to create a cavity., The energy may have been deposited within an existing gap in the H I disk or may have pushed aside the disk material to create a cavity.
+ Alternativelv. the explosion(s) and winds mav have originated in (he nearby. Loop I (Sco-Cen) superbubble and blown into the solar region (Frisch1981:Dochkarev1987).. or the Local Bubble anc Loop I may have been blown independently but are now interacting (Egger&Aschenbach1995)..," Alternatively, the explosion(s) and winds may have originated in the nearby Loop I (Sco-Cen) superbubble and blown into the solar region \citep{frisch_81,bochkarev}, or the Local Bubble and Loop I may have been blown independently but are now interacting \citep{egger_aschenbach}."
+" In some of these scenarios. (he X-ray. emitting plasma may have expanded after being shock heated. vastly cooling the plasma. retarding (he collisional recombination rates. and (hus ""Ireezing the plasma into an ""overionized"" state (Dreitschwerdt&Schumizler1994;Dreitschiwerdt2001)."," In some of these scenarios, the X-ray emitting plasma may have expanded after being shock heated, vastly cooling the plasma, retarding the collisional recombination rates, and thus “freezing” the plasma into an “overionized” state \citep{breitschwerdt_schmutzler,breitschwerdt}."
+. In contrast. if the expansion was restrained and (he time-scale was long. the plasma may have already approached collisional ionizational equilibrium (smith&Cox2001).," In contrast, if the expansion was restrained and the time-scale was long, the plasma may have already approached collisional ionizational equilibrium \citep{smith_cox}."
+. These two cases should produce markedly. different spectral signatures., These two cases should produce markedly different spectral signatures.
+ Presently. many (heorists and observers are working to decipher the Local Dubble's life story by comparing its spectra to spectra predicted by hydrodynamical simulations.," Presently, many theorists and observers are working to decipher the Local Bubble's life story by comparing its X-ray spectra to spectra predicted by hydrodynamical simulations."
+ It is hoped that when the observed fluxes in multiple bands or emission lines are matched by the predictions of a parücular tvpe of model. then we will have found a key to deciphering the story of the Local Bubble.," It is hoped that when the observed fluxes in multiple bands or emission lines are matched by the predictions of a particular type of model, then we will have found a key to deciphering the story of the Local Bubble."
+ Here. we extend (his sort of comparison into the uliraviolet. using the rresonance line (1032. 1038 À)) intensities.," Here, we extend this sort of comparison into the ultraviolet, using the resonance line (1032, 1038 ) intensities."
+ We also compare the intensities with aabsorption line column densities taken lrom the literature. in order to estimate the plvsical properties of the VI--bearmeg plasma.," We also compare the intensities with absorption line column densities taken from the literature, in order to estimate the physical properties of the -bearing plasma."
+" We use (he ""shadowing sirategv long emploved in N-ray. analyses to isolate the Local Dubble's emission from the more distant interstellar mediums emission (see Section 2).", We use the “shadowing” strategy long employed in X-ray analyses to isolate the Local Bubble's emission from the more distant interstellar medium's emission (see Section 2).
+" The observations (228 ksec toward /=278.6"".5 —45.3"") and data reduction techniques are discussed in Section 3."," The observations (228 ksec toward $l = 278.6^{\rm{o}},b = -45.3^{\rm{o}}$ ) and data reduction techniques are discussed in Section 3."
+ Although we observed lor a very long time. no," Although we observed for a very long time, no"
+We have also compared our results with those of other investigations.,We have also compared our results with those of other investigations.
+" Our models made with the most up-to-date input physics agree well with the models of Picardietal. (2004),, Weissetal.(2000) and Siess,Livio&Lattanzio(2002) during evolution on the main sequence and the RGB."," Our models made with the most up-to-date input physics agree well with the models of \citet{pic04}, , \citet{wei00} and \cite{sie02} during evolution on the main sequence and the RGB."
+" In particular, we obtain nearly the same results as do Siess,Livio&Lattanzio(2002),, although we do not follow evolution after the mixing of hydrogen into the helium flash driven convective zone at the beginning of the helium core flash."," In particular, we obtain nearly the same results as do \cite{sie02}, although we do not follow evolution after the mixing of hydrogen into the helium flash driven convective zone at the beginning of the helium core flash."
+" To check the behavior of the He-FDDM event, we need to treat mixing with a time-dependent algorithm."," To check the behavior of the He-FDDM event, we need to treat mixing with a time-dependent algorithm."
+" On the other hand, even after adopting the same radiative opacities, conductive opacities, nuclear reaction rates, and neutrino energy-loss rates as Schlattletal. (2001),, we are not able to obtain the inner ignition at the onset of core helium burning which they find; we suspect that the discrepancy may be due to differences in the EOS for the Coulomb corrections in the liquid and solid states and due to the neglect in their work of neutrino energy losses associated with neutrino bremsstrahlung."," On the other hand, even after adopting the same radiative opacities, conductive opacities, nuclear reaction rates, and neutrino energy-loss rates as \cite{sch01}, we are not able to obtain the inner ignition at the onset of core helium burning which they find; we suspect that the discrepancy may be due to differences in the EOS for the Coulomb corrections in the liquid and solid states and due to the neglect in their work of neutrino energy losses associated with neutrino bremsstrahlung."
+" By treating the resonant and non-resonant rates of reactions separately, we demonstrate that the non-resonant term plays a critical role in the low-mass, zero-metal stars."," By treating the resonant and non-resonant rates of s separately, we demonstrate that the non-resonant term plays a critical role in the low-mass, zero-metal stars."
+" The neglect of non-resonant term causes the lower border in mass of central helium burning, i.e, >1.0Mo."," The neglect of non-resonant term causes the lower border in mass of central helium burning, i.e., $\geq 1.0 \msun$."
+" This explains the discrepancy of the results between ours and Fujimotoetal.(2000), which stems mainly from the difference in the temperature dependence rather than in the energy generation rates themselves."," This explains the discrepancy of the results between ours and \citet{fuj00}, which stems mainly from the difference in the temperature dependence rather than in the energy generation rates themselves."
+ We first point out the possibility of discerning the effect of non-resonant term of ffrom the evolution of stars other than at low temperature regime in the accreting degenerate stars (Nomotoetal.1985)., We first point out the possibility of discerning the effect of non-resonant term of from the evolution of stars other than at low temperature regime in the accreting degenerate stars \citep{nom85}.
+. It is important to precisely determine the abundances and the properties of extremely metal-poor stars to constrain the nuclear reaction rates., It is important to precisely determine the abundances and the properties of extremely metal-poor stars to constrain the nuclear reaction rates.
+ We are grateful to I. Iben Jr. for improving and revising our manuscript., We are grateful to I. Iben Jr. for improving and revising our manuscript.
+ We wish to thank A. Ohnishi and K. Kato for valuable comments on uncertainties on nuclear reaction rates., We wish to thank A. Ohnishi and K. $\bar{\rm o}$ for valuable comments on uncertainties on nuclear reaction rates.
+" This work is part of a PhD. thesis constructed at Hokkaido University and is in part supported by a Grant-in-Aid for Science Research from the Japanese Society for the Promotion ofScience (15204010, 18104003)."," This work is part of a PhD. thesis constructed at Hokkaido University and is in part supported by a Grant-in-Aid for Science Research from the Japanese Society for the Promotion ofScience (15204010, 18104003)."
+in our treatment of the theory of line driven winds. we consistently take into account effects of “uuiltiple-scattering” 1u the transfer of imoinentun frou the radiation field to the wind.,"in our treatment of the theory of line driven winds, we consistently take into account effects of ``multiple-scattering'' in the transfer of momentum from the radiation field to the wind."
+ We find systematic agrecinet between observed and theoretical miass-loss rates for a of O stars., We find systematic agreement between observed and theoretical mass-loss rates for a of O stars.
+ This result iniplies that physical effects that were not incorporated iu our models. such as magnetic fields and stellar rotation. is not expected to influence the mass-loss rates of O stars significantly.," This result implies that physical effects that were not incorporated in our models, such as magnetic fields and stellar rotation, is not expected to influence the mass-loss rates of O stars significantly."
+ lustead of comparing just the imass-loss rates it ds useful to compare (Guodified) wind momenta derived roni observations and theory., Instead of comparing just the mass-loss rates it is useful to compare (modified) wind momenta derived from observations and theory.
+ In earlier studies. c.e. Lamers Leitherer (1993). aud Puls et al. (," In earlier studies, e.g. Lamers Leitherer (1993), and Puls et al. ("
+1996). wind nomenta have been plotted versus the windefficicucy iunnber jg.,"1996), wind momenta have been plotted versus the windefficiency number $\eta$."
+ Comparisons between observed aud theoretical wind momenta as a function of jj could. vicld iuportaut information about the origin of the svstematic discrepancy )etween theory aud observations., Comparisons between observed and theoretical wind momenta as a function of $\eta$ could yield important information about the origin of the systematic discrepancy between theory and observations.
+ Towever. since these two quantities (wind momentum aud wind cficiency nunboer) voth contain the 1iass-loss rate. they are uot independent.," However, since these two quantities (wind momentum and wind efficiency number) both contain the mass-loss rate, they are not independent."
+ Therefore. no such comparison is made here.," Therefore, no such comparison is made here."
+ Iustead. the wind uonieuta are plotted versus the stellar huninosity. ο compare the observational ai theoretical WLR.," Instead, the wind momenta are plotted versus the stellar luminosity, to compare the observational and theoretical WLR."
+ We divido the Tig range into two parts., We divide the $\teff$ range into two parts.
+ First. we exiuunme the wind momenta for stars where Tig= 27 HOO IX. later on we will also compare the cooler stars.," First, we examine the wind momenta for stars where $\teff \ge$ 27 500 K, later on we will also compare the cooler stars."
+ Figure 9 shows the modified wind momenta as a function of stellar huninositv. for he sample of stars with known observational mass-loss rates., Figure \ref{f_modfO} shows the modified wind momentum as a function of stellar luminosity for the sample of stars with known observational mass-loss rates.
+ The upper panel shows these modified wine moment values for the theoretical niss-loss rates and a linear best fit through these theoretical data (dotted line)., The upper panel shows these modified wind momentum values for the theoretical mass-loss rates and a linear best fit through these theoretical data (dotted line).
+" Note that the ""theoretical WLR only coutains the theoretical Af. the iucluded values for ος and & were taken from observations."," Note that the “theoretical” WLR only contains the theoretical $\dot{M}$, the included values for $\vinf$ and $R_*$ were taken from observations."
+ The theoretical WLR is: Since the slope of the WLR of Eq. (153) , The theoretical WLR is: Since the slope of the WLR of Eq. \ref{eq_wlro}) )
+has a slope of w=1.826. the derived theoretical value for eg (Eq. 101) ," has a slope of $x = 1.826$, the derived theoretical value for $\alpha_{\rm eff}$ (Eq. \ref{eq_alphaeff}) )"
+that follows. 1s: This correspouds well to predicted values of the force πιαρα parameter (a2 0.66 and àz 0.10. see e.g. Pauldrach et al.," that follows, is: This corresponds well to predicted values of the force multiplier parameter $\alpha \simeq$ 0.66 and $\delta \simeq$ 0.10, see e.g. Pauldrach et al."
+ 1991)., 1994).
+ The lower pane of Fig., The lower panel of Fig.
+ 9 shows that both the WLR or the Puls et al. (, \ref{f_modfO} shows that both the WLR for the Puls et al. (
+1996) data and that for the other ucthods/authors. follow the same relationship. both iu agreement with the theoretical WLR.,"1996) data and that for the other methods/authors, follow the same relationship, both in agreement with the theoretical WLR."
+ The dotted line is aeain the theoretical best linear fit., The dotted line is again the theoretical best linear fit.
+ We conclude that or the range of the O stars. there is good aerecment vetween theoretical wind momenta aud those determined youn observations.," We conclude that for the range of the O stars, there is good agreement between theoretical wind momenta and those determined from observations."
+ The scatter between theoretical aud observational modified wind momenta is only 0.06 (1 0)., The scatter between theoretical and observational modified wind momenta is only 0.06 (1 $\sigma$ ).
+ The good agreement between he observational aud heoretical wind momenta adds support to the possibility o derive distances to Iuninous. hot stars in extragalactic stellar svsteiis using the WLR.," The good agreement between the observational and theoretical wind momenta adds support to the possibility to derive distances to luminous, hot stars in extragalactic stellar systems using the WLR."
+ Iu oxactiee the technique nay be hampered by e.g. the act that O stars are mostly seen im stellar clusters and cannot be spatially resolved in distant stellar svstems., In practice the technique may be hampered by e.g. the fact that O stars are mostly seen in stellar clusters and cannot be spatially resolved in distant stellar systems.
+ This is one of the reasons why the visually brighter B-type aud especially the A-type superejauts located im the field are expected to be better candidates im actually using the WLR as a distauce indicator (see Iuditzki et al., This is one of the reasons why the visually brighter B-type and especially the A-type supergiants located in the field are expected to be better candidates in actually using the WLR as a distance indicator (see Kudritzki et al.
+ 1999)., 1999).
+ Comparison between the jcoretical aid observational WLR for the winds of B and A type supereiauts is thus essential to investigate wherer the slope of the WLR is the παλιο for different spectral ranges., Comparison between the theoretical and observational WLR for the winds of B and A type supergiants is thus essential to investigate whether the slope of the WLR is the same for different spectral ranges.
+ This is not expected. since the winds of different spectral types are driven by lines of cifferent ious (see Vink et al.," This is not expected, since the winds of different spectral types are driven by lines of different ions (see Vink et al."
+ 1999: Puls et al., 1999; Puls et al.
+ 2000)., 2000).
+ Figure 10 shows the modified wind momentum as a function of luminosity for both theory aud observations for the stars iu the second range (12 500 <Έως 22 500 Kk)., Figure \ref{f_modfB} shows the modified wind momentum as a function of luminosity for both theory and observations for the stars in the second range (12 500 $\le \teff \le$ 22 500 K).
+ A best fit through the theoretically derived WLR is indicated with a dotted line iu both pancls., A best fit through the theoretically derived WLR is indicated with a dotted line in both panels.
+ The theoretical WLR for this Tig range is: Since the slope of the WLR for this range is slightly higher than that for the O star range. the predicted value for ag is somewhat lower (see Sect. 1.3)).," The theoretical WLR for this $\teff$ range is: Since the slope of the WLR for this range is slightly higher than that for the O star range, the predicted value for $\alpha_{\rm eff}$ is somewhat lower (see Sect. \ref{s_mwm}) ),"
+ uamely: The lower panel of Fig., namely: The lower panel of Fig.
+ 10 indicates the observed modified wind momenta (the dotted line contaius the theoretical wiass-loss rates)., \ref{f_modfB} indicates the observed modified wind momenta (the dotted line contains the mass-loss rates).
+ For this secoud Tig rauge (42 500 €T4< 22 500 Is) the plot in the lower paucl reveals a large scatter in the observed data., For this second $\teff$ range (12 500 $\le \teff \le$ 22 500 K) the plot in the lower panel reveals a large scatter in the observed data.
+The ability to determine the ages of. pre-main-sequence (PAIS) stars is crucial for advancing our understanding of the carly phases of stellar evolution.,The ability to determine the ages of pre-main-sequence (PMS) stars is crucial for advancing our understanding of the early phases of stellar evolution.
+ Phere are. two kev applications., There are two key applications.
+ Firsth. and perhaps most obviously. we need stellar ages if one is to carry our experiments such as tracing the evolution of stellar angular momentum. or ollowing the fraction of stars with proto-planctary disces as a Function of time.," Firstly, and perhaps most obviously, we need stellar ages if one is to carry our experiments such as tracing the evolution of stellar angular momentum, or following the fraction of stars with proto-planetary discs as a function of time."
+ Secondly. for PAIS stars. the conversion rom observables such as temperature and luminosity into mass is highly age dependent. making accurate ages vita or determining the mass unction.," Secondly, for PMS stars, the conversion from observables such as temperature and luminosity into mass is highly age dependent, making accurate ages vital for determining the mass function."
+ The primary method of determining the ages required for these studies is to compare he observed properties of PAIS stars with mocels., The primary method of determining the ages required for these studies is to compare the observed properties of PMS stars with models.
+ Phe mos easily. accessible obscrvables are a star's temperature an uminosityv. since they can be measured [rom its colours anc magnitudes.," The most easily accessible observables are a star's temperature and luminosity, since they can be measured from its colours and magnitudes."
+ Phe problem is that for the same colours ane magnitudes cifferent models can predict ages which cliller w a factor two. and. even the same models will predic different ages depending on which colours ancl magnitudes are used.," The problem is that for the same colours and magnitudes different models can predict ages which differ by a factor two, and even the same models will predict different ages depending on which colours and magnitudes are used."
+ This makes meaningful comparisons between the ages quoted in the literature for clusters or associations a best dillicult. ancl often impossible.," This makes meaningful comparisons between the ages quoted in the literature for clusters or associations at best difficult, and often impossible."
+ It was these problems which led. us to devise a mocel-independent. age ordering of voung clusters and associations based. on their diagrams (τι., It was these problems which led us to devise a model-independent age ordering of young clusters and associations based on their diagrams \citep{2007MNRAS.375.1220M}.
+ For PAIS stars the primary age diagnostic is based on the fact that stars [ade as they get older ancl contract tow:uds the main sequence (ALS)., For PMS stars the primary age diagnostic is based on the fact that stars fade as they get older and contract towards the main sequence (MS).
+ We used this movement of the xquence towards progressively fainter magnitudes to derive an age ordering. although to do so we also had to measure! à Consistent. set of clistances. which we derived [rom he niore massive stars which have already. reached the MS (?)," We used this movement of the sequence towards progressively fainter magnitudes to derive an age ordering, although to do so we also had to measure a consistent set of distances, which we derived from the more massive stars which have already reached the MS \citep{2008MNRAS.386..261M}."
+ Whilst an age ordering such as ours is useful. for example it has showed unambiguously that different clusters take cdillerent times to reach the same cise fraction or angular momentum clistribution. for quantitative work an absolute scale is required.," Whilst an age ordering such as ours is useful, for example it has showed unambiguously that different clusters take different times to reach the same disc fraction or angular momentum distribution, for quantitative work an absolute scale is required."
+ For PAIS clusters ancl associations were are several usable age indicators. each of which relics n Comparing stellar properticSs with mocels.," For PMS clusters and associations there are several usable age indicators, each of which relies on comparing stellar properties with models."
+ For this reason is best to group them accorcling to the underlying physics., For this reason it is best to group them according to the underlying physics.
+ =μι is the contraction of PAIS stars as they approach ve MS., First is the contraction of PMS stars as they approach the MS.
+ As pointed out above. and. discussed at length in these “contraction” or “PAIS” ages are highly. model ependent. and given the euwrent disagreements between 10 models cannot νο] an absolute age scale.," As pointed out above, and discussed at length in \cite{2007MNRAS.375.1220M} these “contraction” or “PMS” ages are highly model dependent, and given the current disagreements between the models cannot yield an absolute age scale."
+ Although most stars in a voung cluster or association are in the PAIS hase. the evolution of the most massive stars proceeds SO fast. that they may not only have reached the AIS. out. evolved: ονομα it.," Although most stars in a young cluster or association are in the PMS phase, the evolution of the most massive stars proceeds so fast that they may not only have reached the MS, but evolved beyond it."
+ This gives us access to (wo more age measures., This gives us access to two more age measures.
+ First. having reached the MS. stars move," First, having reached the MS, stars move"
+The formation of the first barvonic objects (in particular the first stars) was an important. milestone in the history of the Universe.,The formation of the first baryonic objects (in particular the first stars) was an important milestone in the history of the Universe.
+" lt marked the transition between the cold. neutral. metal-free universe (the epoch called the ""dark ~ cosmological ages that started. right after recombination) and the mocerm ionized. hot. anc metal-rich universe."," It marked the transition between the cold, neutral, metal-free universe (the epoch called the “dark ” cosmological ages that started right after recombination) and the modern ionized, hot, and metal-rich universe."
+ The ormation of the very first stars is expected to be relatively simple: this is due to the primordial chemistry before stars oduced heavy. elements. and the simplified gas dynamics in the absence of cynamically-relevant magnetic fields and οσοασ from. luminous objects (Leemarketal.1997:Darkana&Loeb2001 ).," The formation of the very first stars is expected to be relatively simple; this is due to the primordial chemistry before stars produced heavy elements, and the simplified gas dynamics in the absence of dynamically-relevant magnetic fields and feedback from luminous objects \citep{Tegmark:1997,Barkana:2001}."
+. Since molecular hydrogen line emission is the lowest-emperature coolant in metal-free gas. the first stars are expected to have formed in. halos with total mass above ~10° AL. (Loemarketal.1997).," Since molecular hydrogen line emission is the lowest-temperature coolant in metal-free gas, the first stars are expected to have formed in halos with total mass above $\sim 10^5$ $_\odot$ \citep{Tegmark:1997}."
+. More generally. if the mass ofa dark matter halo is higher than a threshold referred to as the minimum cooling mass CM). the collapsing gas is heated to a high enough temperature that it emits radiation.," More generally, if the mass of a dark matter halo is higher than a threshold referred to as the minimum cooling mass $M\cool$ ), the collapsing gas is heated to a high enough temperature that it emits radiation."
+ lt then cools and condenses. allowing a star to form.," It then cools and condenses, allowing a star to form."
+" The threshold. can also be described. as a minimum circular velocity. (V5,4) via the standard relation V=Y/CAL/R| for a halo of mass AZ and virial radius 2."," The threshold can also be described as a minimum circular velocity $V\cool$ ) via the standard relation $V_c =
+\sqrt{GM/R}$ for a halo of mass $M$ and virial radius $R$."
+ This scenario of the earliest star formation has been confirmed by numerical simulations using both Adaptive Mesh Refinement (AATR) and. Smooth Particle Lvedrocdyvnamics (SPLHD codes (e.g...Puller&Couchman2000:al.2006:Turkοἱ 2011).," This scenario of the earliest star formation has been confirmed by numerical simulations using both Adaptive Mesh Refinement (AMR) and Smooth Particle Hydrodynamics (SPH) codes \citep[e.g.,][]{Fuller:2000,
+Abel:2001, Bromm:2002, Yoshida:2003, Reed:2005, Yoshida:2006,
+Turk:2011}."
+.. All these simulations did not account for the initial relative velocities between the barvons and the dark matter., All these simulations did not account for the initial relative velocities between the baryons and the dark matter.
+ We now summarize two recent SPLI simulations (Stacyetal.2010:Cireif2011) that studied the impact of the relative streaming velocity en. on the formation of the first stars.," We now summarize two recent SPH simulations \citep{Stacy:2011, Greif:2011}
+ that studied the impact of the relative streaming velocity $v\bc$ on the formation of the first stars."
+ Numerical simulations face a great cilliculty at high redshift. since they must. resolve. the then-tvpical tiny galaxies while at the same time capture the global galaxy distribution which is characterized. by strong [luctuations on surprisingly large scales (Barkana&Loeb2004).," Numerical simulations face a great difficulty at high redshift, since they must resolve the then-typical tiny galaxies while at the same time capture the global galaxy distribution which is characterized by strong fluctuations on surprisingly large scales \citep{Barkana:2004}."
+. The relative velocities are correlated up to scales above 100 Alpe. and they are important at. high redshifts where star formation is dominated by very small halos.," The relative velocities are correlated up to scales above 100 Mpc, and they are important at high redshifts where star formation is dominated by very small halos."
+ Cosmological simulations that cover this range of scales are not currently feasible., Cosmological simulations that cover this range of scales are not currently feasible.
+ Llowever. numerical simulations are the best tool for studying the complex. non-linear formation of halos on small scales.," However, numerical simulations are the best tool for studying the complex, non-linear formation of halos on small scales."
+ The scales relevant to the formation of the small halos that host the first stars are well below the coherence scale of the relative velocity field., The scales relevant to the formation of the small halos that host the first stars are well below the coherence scale of the relative velocity field.
+ Pherefore it is possible to simulate halo formation in small patches of uniform ei., Therefore it is possible to simulate halo formation in small patches of uniform $v\bc$.
+ The simulations vield the mass reached. by a halo when it first allows a star to form. i.c. when it first contains a cooling. rapidly-collapsing gas core.," The simulations yield the mass reached by a halo when it first allows a star to form, i.e., when it first contains a cooling, rapidly-collapsing gas core."
+ The results show a substantially increased halo mass in regions with a significant relative velocity., The results show a substantially increased halo mass in regions with a significant relative velocity.
+ This is a different effect. from the suppression of the amount of eas. which implies a smaller number of stars in the halo at a given time: instead in this case there is a substantial celay in the formation of the first star within the halo.," This is a different effect from the suppression of the amount of gas, which implies a smaller number of stars in the halo at a given time; instead in this case there is a substantial delay in the formation of the first star within the halo."
+ Moreover. this effect is not simply related to the total amount of accreted gas. since in the cases with a bulk velocity. even if we wait for the halo to accrete the same total gas mass as its no-velocity counterpart. it still does not form a star (even within the now deeper potential of a more massive host halo): the delay. is substantially longer than would. be expected: based on a fixed. total mass of accreted gas.," Moreover, this effect is not simply related to the total amount of accreted gas, since in the cases with a bulk velocity, even if we wait for the halo to accrete the same total gas mass as its no-velocity counterpart, it still does not form a star (even within the now deeper potential of a more massive host halo); the delay is substantially longer than would be expected based on a fixed total mass of accreted gas."
+ Instead. it appears that the explanation Lies with the internal density and temperature profiles of the gas. which are strongly alleeted by the presence of the streaming motion.," Instead, it appears that the explanation lies with the internal density and temperature profiles of the gas, which are strongly affected by the presence of the streaming motion."
+ A plausible explanation for the resulting delay in star formation is that the first star forms from the gas that would have accreted early and formed the dense central cores in which stars form: this gas tries to accrete early (when vei is still very large) into a still-nall halo progenitor. so it is πα most. strongly by the suppression of gas accretion due to the bulk velocity.," A plausible explanation for the resulting delay in star formation is that the first star forms from the gas that would have accreted early and formed the dense central cores in which stars form; this gas tries to accrete early (when $v\bc$ is still very large) into a still-small halo progenitor, so it is affected most strongly by the suppression of gas accretion due to the bulk velocity."
+ The simulations vield a minimum halo cooling mass at various redshifts. so we fit the results to find the dependence of the minimum halo mass on the redshift of collapse and on the bulk velocity. enc. in the patch.," The simulations yield a minimum halo cooling mass at various redshifts, so we fit the results to find the dependence of the minimum halo mass on the redshift of collapse and on the bulk velocity, $v\bc$, in the patch."
+ This will then allow us to study the ellect of the relative velocity on the formation of the first stars using statistical methods that average over large cosmological regions that cannot be directly simulated., This will then allow us to study the effect of the relative velocity on the formation of the first stars using statistical methods that average over large cosmological regions that cannot be directly simulated.
+ Stacyctal.(2010) ane Greifetal.(2011) state apparently contradictory. conclusions. one claiming a =-egligible elfect on star-forming halos and the other a large feet.," \citet{Stacy:2011} and \citet{Greif:2011} state apparently contradictory conclusions, one claiming a negligible effect on star-forming halos and the other a large effect."
+ In order to meaningfully compare their results. it is important to put them both on the same scale.," In order to meaningfully compare their results, it is important to put them both on the same scale."
+ We express the cooling threshold. as a halo circular. velocity. since simulations (cited above) without the bulk velocity find an approximately redshift-independent: threshold. Vino: this is naturally expected since molecular cooling turns on essentially at a fixed gas temperature. and the halo circular velocity. determines the virial temperature to which the eas is heated.," We express the cooling threshold as a halo circular velocity, since simulations (cited above) without the bulk velocity find an approximately redshift-independent threshold $V\coolz $; this is naturally expected since molecular cooling turns on essentially at a fixed gas temperature, and the halo circular velocity determines the virial temperature to which the gas is heated."
+ Thus. the limit of zero bulk velocity simply eives a fixed threshold. Vine.," Thus, the limit of zero bulk velocity simply gives a fixed threshold $V\coolz $."
+ When we add the relative velocities. in. principle the minimum circular velocity in a patch may. be a separate function of two parameters. the redshift z and the bulk velocity at halo formation e«(z).," When we add the relative velocities, in principle the minimum circular velocity in a patch may be a separate function of two parameters, the redshift $z$ and the bulk velocity at halo formation $v\bc(z)$."
+" Vhe history of ci; at earlier. redshifts cannot. introduce aclelitional parameters. since given. both z and me(s). the full history of ey. is determined. Le. at any other recshift Sorel)Sele).(1fC12),"," The history of $v\bc$ at earlier redshifts cannot introduce additional parameters, since given both $z$ and $v\bc(z)$, the full history of $v\bc$ is determined, i.e., at any other redshift $z'$, $v\bc(z')=v\bc(z) \times
+(1+z')/(1+z)$."
+ Consider now the limit of a very high. bulk velocity. mele)2Ἑλωμηνον so that the clleet of Vossio is negligible.," Consider now the limit of a very high bulk velocity, $v\bc(z) \gg
+V\coolz $, so that the effect of $V\coolz $ is negligible."
+ For simplicity. consider for a moment a constant ey. versus redshift. fixed at its final value μις) at the halo formation redshift z.," For simplicity, consider for a moment a constant $v\bc$ versus redshift, fixed at its final value $v\bc(z)$ at the halo formation redshift $z$."
+ In this case there is only one velocity scale in the problem., In this case there is only one velocity scale in the problem.
+ As ina Jeans mass analysis. in the reference frame of a collapsing dark matter halo with a circular velocity ὃν. clearly gravity will be able to pull in the gas (which streams by at the velocity euo(2)) iV esos).," As in a Jeans mass analysis, in the reference frame of a collapsing dark matter halo with a circular velocity $V_c$, clearly gravity will be able to pull in the gas (which streams by at the velocity $v\bc(z)$ ) if $V_c \ga v\bc(z)$ ."
+ Now. in the real case where e) is higher during the formation of the halo. we would expect to get a threshold that is higher than eic(z). but by a fixed factor. because the physics is scale-free: on one side. ei scales in a simple way with redshift. and on the other side. halo formation (in the high-redshift. Einstein universe) also scales in a simple way. às we know from spherical collapse: e.g.. turnaround for a halo that formis at redshiltz always occurs at z/ where 1|c/1.59(12) so that ei«(z/)= 1.5960(2).," Now, in the real case where $v\bc(z')$ is higher during the formation of the halo, we would expect to get a threshold that is higher than $v\bc(z)$, but by a fixed factor, because the physics is scale-free: on one side, $v\bc$ scales in a simple way with redshift, and on the other side, halo formation (in the high-redshift, Einstein de-Sitter universe) also scales in a simple way, as we know from spherical collapse; e.g., turnaround for a halo that forms at redshift$z$ always occurs at $z'$ where $1+z' = 1.59 (1+z)$ so that $v\bc(z') = 1.59 v\bc(z)$ ."
+ Phe only new scale that enters, The only new scale that enters
+the density region higher than ~1015 g/cm?.,the density region higher than $\sim 10^{13}$ $^3$.
+ Figure 3 shows the unstable region for a relativistic Hartree mean field model in which the baryon effective mass is taken into account., Figure 3 shows the unstable region for a relativistic Hartree mean field model in which the baryon effective mass is taken into account.
+" When we consider the effective baryon mass within the relativistic Hartree theory, domain formation can significantly occur above a density of p1014 g/cm? and field strength B~1016 G. The effective baryon mass alowers the density at which magnetic domain formation occurs in the core of a magnetar in which strong magnetic fields are expected."," When we consider the effective baryon mass within the relativistic Hartree theory, domain formation can significantly occur above a density of $\rho \sim 10^{14}$ $^{3}$ and a field strength $B \sim 10^{16}$ G. The effective baryon mass lowers the density at which magnetic domain formation occurs in the core of a magnetar in which strong magnetic fields are expected."
+" We also find that magnetic domain formation could not be formed for a magnetic field strength of <1015 G. Finally, in Figure 4, the magnetic domain instability regions obtained by using the magnetic BPS model are depicted in the outer crust of magnetars."," We also find that magnetic domain formation could not be formed for a magnetic field strength of $\lesssim 10^{18}$ G. Finally, in Figure 4, the magnetic domain instability regions obtained by using the magnetic BPS model are depicted in the outer crust of magnetars."
+" Unlike the other two cases, in this case the magnetization is dominated by electrons."," Unlike the other two cases, in this case the magnetization is dominated by electrons."
+ Therefore the conditions for the onset of the magnetic domain instability is affected by the occupation of Landau levels., Therefore the conditions for the onset of the magnetic domain instability is affected by the occupation of Landau levels.
+ This causes the instability conditions to vary as a function of field strength for fixed density., This causes the instability conditions to vary as a function of field strength for fixed density.
+" In Figure 4 we see that magnetic domains cannot be formed if the density is less than p~5x1010 g/cm? for a magnetar having a typical surface magnetic field of B~1014-15 G. However, in the density region around the neutron drip density, magnetic domains can be formed."," In Figure 4 we see that magnetic domains cannot be formed if the density is less than $\rho \sim 5 \times 10^{10}$ ${^3}$ for a magnetar having a typical surface magnetic field of $B \sim 10^{14 - 15}$ G. However, in the density region around the neutron drip density, magnetic domains can be formed."
+ We can also see magnetic domain formation in regions at low density >1011 g/cm? and at a relatively low magnetic field strength >10!? G. This is the case considered by Blandford and Hernquist (1982)., We can also see magnetic domain formation in regions at low density $\gsim 10^{11}$ $^3$ and at a relatively low magnetic field strength $\gsim 10^{13}$ G. This is the case considered by Blandford and Hernquist (1982).
+" For magnetic white dwarfs, Adam (1986) also obtained the similar result for a magnetized electron gas."," For magnetic white dwarfs, Adam (1986) also obtained the similar result for a magnetized electron gas."
+" We thus find that possible unstable regions for magnetic domain formation are within the deep outer crust, deeper part of the inner crust, and core of magnetars."," We thus find that possible unstable regions for magnetic domain formation are within the deep outer crust, deeper part of the inner crust, and core of magnetars."
+ This means that the outer shell at low density less than pc1010 g/cm? is stable against magnetic domain formation and should consists of strongly magnetized material without magnetic domains., This means that the outer shell at low density less than $\rho \sim 10^{10}$ ${^3}$ is stable against magnetic domain formation and should consists of strongly magnetized material without magnetic domains.
+ This magnetic domain formation might be an important clue to explain SGRs and AXPs in the magnetar model., This magnetic domain formation might be an important clue to explain SGRs and AXPs in the magnetar model.
+" As the density increases, the number of Landau orbital also increases."," As the density increases, the number of Landau orbital also increases."
+" Therefore, Based upon Blandford & Hermquist we can calculate the thickness of the crust, Az, (1982),using the gravitational potential energy of a nucleon and the Fermi energy εε per electron."," Therefore, Based upon Blandford $\&$ Hermquist (1982), we can calculate the thickness of the crust, $\Delta z$, using the gravitational potential energy of a nucleon and the Fermi energy $\epe$ per electron."
+ We also know the maximum number of the Landau orbitals n=(c2— where γε=B/B£.," We also know the maximum number of the Landau orbitals $n = (\epe^2 - 1) / 2 \gme$, where $\gme = B / B_{c}^{e}$."
+" For a fixed B, An~εεΔεε/Β. 1)/25,"," For a fixed B, $\Delta n \sim \epe \Delta \epe / B$."
+",Therefore, we can get Az/An."," Therefore, we can get $\Delta z / \Delta n$."
+" For An—1, we finally can estimate the spacing between layers associated with the maximum Landau level: Az 100 (sis), where g is the οsurface gravity, µε is (e)the mean («)molecular weight per electron, and e,=E%/m,-c? denotes the electron Fermi energy."," For $\Delta n = 1$, we finally can estimate the spacing between layers associated with the maximum Landau level: z 100 ( ( ), where $g$ is the surface gravity, $\mu_e$ is the mean molecular weight per electron, and $\epe = E_{F}^{e} /m_e c^2$ denotes the electron Fermi energy."
+" Under the conditions for domain formation, the spacing Az would be the average vertical scale between domain interfaces as well as a horizontal scale of the domains if they are in local pressure"," Under the conditions for domain formation, the spacing $\Delta z$ would be the average vertical scale between domain interfaces as well as a horizontal scale of the domains if they are in local pressure"
+given by where Go is the center-surface density contrast at the time that the molecular cloud turbulence was inherited.,given by where $\zeta_0$ is the center–surface density contrast at the time that the molecular cloud turbulence was inherited.
+ The spectrum of supersonic turbulence in larger scale molecular clouds is not well understood., The spectrum of supersonic turbulence in larger scale molecular clouds is not well understood.
+" Observations indicate that it is close to Kolmogorov, in the range of o~L!? to c~L1? where c is the observed velocity dispersion and L the length scale."," Observations indicate that it is close to Kolmogorov, in the range of $\sigma \sim L^{-1/3}$ to $\sigma \sim L^{-1/2}$ where $\sigma$ is the observed velocity dispersion and $L$ the length scale."
+" For consistency with our previous analysis, we will consider both our Kolmogorov and Flat energy spectra as the inherited spectra."," For consistency with our previous analysis, we will consider both our Kolmogorov and Flat energy spectra as the inherited spectra."
+" For the same reason that the fundamental modes (η= 0) and p-modes (n> 0) produce different contributions to Pss, an adiabatic contraction results in a energy spectrum if Co is sufficiently small."," For the same reason that the fundamental modes $n=0$ ) and p-modes $n>0$ ) produce different contributions to $\Po$, an adiabatic contraction results in a fundamental-heavy energy spectrum if $\zeta_0$ is sufficiently small."
+fundamental-heavy This is shown at ¢=14 in Figure 6 for both the Kolmogorov and Flat energy spectra with Go=1.25., This is shown at $\zeta=14$ in Figure \ref{fig:Eevol} for both the Kolmogorov and Flat energy spectra with $\zeta_0=1.25$.
+" During the collapse phase, Ε.ΟΕΦ. has decreased from 3.5 and 3.9 to 1.3 and 1.5 for the Kolmogorov and Flat energy spectra, respectively."," During the collapse phase, $\Eo^{>0}/\Eo^0$ has decreased from $3.5$ and $3.9$ to $1.3$ and $1.5$ for the Kolmogorov and Flat energy spectra, respectively."
+" Thus, while in neither case the fundamental modes dominate, in both the energy in the fundamentals and p-modes are comparable."," Thus, while in neither case the fundamental modes dominate, in both the energy in the fundamentals and p-modes are comparable."
+" This suggests that neither EE, nor EK may be applicable, but rather some combination of the two is appropriate for the observed pulsating starless cores."," This suggests that neither $E_{nlm}^{K}$ nor $E_{nlm}^{fK}$ may be applicable, but rather some combination of the two is appropriate for the observed pulsating starless cores."
+" This also provides a natural for the dominance of low-order modes, and in particular explanationlow-order quadrupole modes, as implicated by self-absorbed molecular line spectra in Barnard 68 (Ladaetal.asymmetric2003;Keto2006;Redmanericketal. 2007)."," This also provides a natural explanation for the dominance of low-order modes, and in particular low-order quadrupole modes, as implicated by self-absorbed asymmetric molecular line spectra in Barnard 68 \citep{Lada2003,Keto2006,Redman2006,Broderick2007}."
+. In Section ?? we showed how pulsations can super-critical isothermal configurations., In Section \ref{sec:AME} we showed how pulsations can support super-critical isothermal configurations.
+" However, in the supportabsence of continual excitation, these oscillations eventually "," However, in the absence of continual excitation, these oscillations eventually decay."
+The manner and timescale in which large-amplitude decay.pulsations damp depends upon the underlying core parameters and environment as well as the detailed structure of each individual oscillation and their amplitudes., The manner and timescale in which large-amplitude pulsations damp depends upon the underlying core parameters and environment as well as the detailed structure of each individual oscillation and their amplitudes.
+ In general the lifetimes are different for each mode., In general the lifetimes are different for each mode.
+ In our previous papers we have estimated life times for a few particular cases., In our previous papers we have estimated life times for a few particular cases.
+ Figures 2 and 5 of Brodericketal.(2007) estimate a decay time of 2 Myr for the 122 and 111 modes by mode-mode coupling., Figures 2 and 5 of \citet{Broderick2007} estimate a decay time of 2 Myr for the 122 and 111 modes by mode-mode coupling.
+" In Brodericketal. (2008),, we estimate decay times of 1 and 0.4 Myr by transmission to the larger scale ISM for Myrisolated and embedded cores "," In \citet{Broderick2008}, we estimate decay times of 1 Myr and 0.4 Myr by transmission to the larger scale ISM for isolated and embedded cores respectively."
+"These examples are too few to define the decay times respectively.precisely; therefore, for simplicity, we assume that all pulsations decay exponentially with the same lifetime, 7~106yr, chosen to be consistent with our previous numerical simulations of the non-linear mode-mode and environment couplings (Brodericketal.2007,2008)."," These examples are too few to define the decay times precisely; therefore, for simplicity, we assume that all pulsations decay exponentially with the same lifetime, $\tau\simeq 10^6\,\yr$, chosen to be consistent with our previous numerical simulations of the non-linear mode-mode and core-environment couplings \citep{Broderick2007,Broderick2008}."
+. The consequences of mode damping are shown in Figure 7.., The consequences of mode damping are shown explicitly in Figure \ref{fig:evol}.
+" Here P(R) curves are shown for a number of explicitlymode where the is that resulting from a slow amplitudes,initial contraction energyperiod, spectrumdescribed in Section 4.2.."," Here $P(R)$ curves are shown for a number of mode amplitudes, where the energy spectrum is that resulting from a slow initial contraction period, described in Section \ref{sec:ESfRT}."
+" these are normalized such that of the binding energy Initially,is in oscillations at ¢=14 (dark-blue line), and then at t/r=0.5, 1.0, and 1.5."," Initially, these are normalized such that of the binding energy is in oscillations at $\zeta=14$ (dark-blue line), and then at $t/\tau=0.5$, $1.0$, and $1.5$."
+" Since the exterior is constant, we show a number of points with a fixed, super-criticalpressure surface pressure, illustrating the onset of instability during Pulsation-decay driven evolution."," Since the exterior pressure is constant, we show a number of points with a fixed, super-critical surface pressure, illustrating the onset of instability during Pulsation-decay driven evolution."
+" The associated evolution in ¢, which in this case is simply proportional to the central density, is shown in the inset, with the colored points corresponding to the curves in the main Figure."," The associated evolution in $\zeta$, which in this case is simply proportional to the central density, is shown in the inset, with the colored points corresponding to the curves in the main Figure."
+" Prior to the times shown, we assume the cores were forming with T<r, presumably during a phase of prolonged contraction?."," Prior to the times shown, we assume the cores were forming with $T<\tau$, presumably during a phase of prolonged contraction."
+". At some point, due to both the increase in the dynamical importance of the pulsations and the decrease in T due to the increased mode amplitudes and central density, the oscillations begin to dominate the formation processes, and the cores enter into a pulsation-driven phase."," At some point, due to both the increase in the dynamical importance of the pulsations and the decrease in $\tau$ due to the increased mode amplitudes and central density, the oscillations begin to dominate the formation processes, and the cores enter into a pulsation-driven phase."
+" During this time cores can derive a substantial portion of their pressure from the oscillations themselves, and their damping drives supportchanges in the underlying configurations on timescales comparable to 7."," During this time cores can derive a substantial portion of their pressure support from the oscillations themselves, and their damping drives changes in the underlying configurations on timescales comparable to $\tau$."
+" Finally, for super-critical cores, once the energy remaining in pulsations has reached some critical value, oscillations can no longer support the core, leading to unstable collapse."," Finally, for super-critical cores, once the energy remaining in pulsations has reached some critical value, oscillations can no longer support the core, leading to unstable collapse."
+" Thus, we may think of the evolution of a pulsating starless core as consisting of at least three phases: Formation, Pulsation-supported and Collapse."," Thus, we may think of the evolution of a pulsating starless core as consisting of at least three phases: Formation, Pulsation-supported and Collapse."
+ We have already used the existence of large-amplitude oscillations to empirically determine the relationship between the mode excitation timescale and the other relevant timescales., We have already used the existence of large-amplitude oscillations to empirically determine the relationship between the mode excitation timescale and the other relevant timescales.
+" However, these timescales are not generally independent of the"," However, these timescales are not generally independent of the"
+2INDe (TT).," \citep{ww95,nomoto-etal97}."
+ Therefore. a top-leavy IMF may produce more à elemieuts than a population with a normal IME.," Therefore, a top-heavy IMF may produce more $\alpha$ elements than a population with a normal IMF."
+ Cirrent Galactic chemical evolution models accurately reproduce the observed abundance patterns of Galactic stars over a large rauge of metalliities. leaving Little rooni for siguificant changes.," Current Galactic chemical evolution models accurately reproduce the observed abundance patterns of Galactic stars over a large range of metallicities, leaving little room for significant changes."
+ However. oue clement whose abuudances are not well reproduced is Ale. which is systematically underpredieted by the models over the range 2.5€οX0.5 (andslieltlyover-predictedatlowerietallicities:seefigure1of 7).," However, one element whose abundances are not well reproduced is Mg, which is systematically underpredicted by the models over the range $-2.5 \la \overH{Fe} \la -0.5$ \citep[and slightly overpredicted at lower
+metallicities; see figure~1 of][]{francois-etal04}."
+ Tuterestinely. these are precisely the metallicities where most of the CAIB+ceulated star formation occurs within our simmlatious. aud would indeed be produced at ereater rates with a top-heavy IME.," Interestingly, these are precisely the metallicities where most of the CMB-regulated star formation occurs within our simulations, and would indeed be produced at greater rates with a top-heavy IMF."
+ However. the QO abundances do not support this conjecture: they are. if anvthine at these metallicities. while a top-heavy IMF would have au even larecr effect ou than on2!Mg.," However, the O abundances do not support this conjecture: they are, if anything at these metallicities, while a top-heavy IMF would have an even larger effect on than on."
+ Uuless some uiaccouuted-for process svsteniaticallv drives down the abundance by a sienificant amount. a lore parsimomous explanation is that the vields require adjustineut. as suggested by ?..," Unless some unaccounted-for process systematically drives down the abundance by a significant amount, a more parsimonious explanation is that the yields require adjustment, as suggested by \citet{francois-etal04}."
+ We have investigated the sugeestion of 7. and ? that star formation at high redshift aud high metallicity results in a top-heavy IME due to the influence of the CNID as an effective temperature floor for the collapsing protostellay core., We have investigated the suggestion of \citet{larson05} and \citet{stson} that star formation at high redshift and high metallicity results in a top-heavy IMF due to the influence of the CMB as an effective temperature floor for the collapsing protostellar core.
+ By asstuuing that CMD reeulation is iuportant if the mininuun temperature the core would reach in the absence of the CAIB. according to the models of 7.. as less than the CAIB temperature at a given redshift. we are able to parametrize the evolution of the critical metallicity aas a function of redshift.," By assuming that CMB regulation is important if the minimum temperature the core would reach in the absence of the CMB, according to the models of \citet{omukai-etal05}, is less than the CMB temperature at a given redshift, we are able to parametrize the evolution of the critical metallicity as a function of redshift."
+ Comparison with the 3D lydrodvuamiucal simulations of ? sugeests that our model is generally accurate. but may underestimate bby up to an order of magnitude in particular cases.," Comparison with the 3D hydrodynamical simulations of \citet{stson} suggests that our model is generally accurate, but may underestimate by up to an order of magnitude in particular cases."
+ By comparing tto the metallicities of stays formed in high resolution cosinological simulatious of the formation of cight ~L. galaxies. we conclude that if CMD regulation. does operate. the majority of stars at redshifts between 2.7 and 6 would have formed in the CMD-regulated reeinue.," By comparing to the metallicities of stars formed in high resolution cosmological simulations of the formation of eight $\sim L_*$ galaxies, we conclude that if CMB regulation does operate, the majority of stars at redshifts between $2.7$ and $6$ would have formed in the CMB-regulated regime."
+ We have investigated five possible observational sicnatures of CMD regulation: Iu conclusion. the observatious neither provide conclusive evidence for the hwpothesis that the CMD can regulate star formation. nor rule it out.," We have investigated five possible observational signatures of CMB regulation: In conclusion, the observations neither provide conclusive evidence for the hypothesis that the CMB can regulate star formation, nor rule it out."
+ The most convincing evidence would come iu the form of either the clear presence or absence of an envelope in redshift-ietallicitv space. above which stars are either abseut or have ereatly reduced iunibers compared to the number of stars believed. to have formed there from other ΠΠΠο uufortuuatelv. as discussed in L1. it is unclear if stellar ages will ever be known preciscly enough to perform this test.," The most convincing evidence would come in the form of either the clear presence or absence of an envelope in redshift-metallicity space, above which stars are either absent or have greatly reduced numbers compared to the number of stars believed to have formed there from other measurements; unfortunately, as discussed in \ref{section:gcs}, it is unclear if stellar ages will ever be known precisely enough to perform this test."
+ Another potential avenue of exploration would be to use the methods that currently arene for a imean chanec in the IMPE with time. such as in ? and 7.. but subdividiug the populations by metallicity o determine if au evolving IME is only required for uetalaieh stars. as would be expected if the evolution iu the IMPE is due to CMD regulation. or if it must be τιν universal.," Another potential avenue of exploration would be to use the methods that currently argue for a mean change in the IMF with time, such as in \citet{dave08} and \citet{wth08}, but subdividing the populations by metallicity to determine if an evolving IMF is only required for metal-rich stars, as would be expected if the evolution in the IMF is due to CMB regulation, or if it must be truly universal."
+ This paper makes use of simulations performed as mart of the SIITARCNET Dedicated. Resource project: UMUCS: The MeMaster Unbiased. Galaxy. Simulations Project” (DR316. DRLOL. aud DR137).," This paper makes use of simulations performed as part of the SHARCNET Dedicated Resource project: “MUGS: The McMaster Unbiased Galaxy Simulations Project” (DR316, DR401, and DR437)."
+Here. as belore. while r; is now defined in terms of theΜα mass:,"Here, as before, while $r_s$ is now defined in terms of the mass:."
+ The fully nondimenusional evolutionary equation for the speed is then In this last equation. we have introduced the initial. nondiniensioual speed η. as well as a nondimensional time.//ly.. where The denominator in equation (73)) is the filucial mass accretion rate defiued in Section L1.," The fully nondimensional evolutionary equation for the speed is then In this last equation, we have introduced the initial, nondimensional speed $\beta_0$, as well as a nondimensional time, where The denominator in equation \ref{eqn:veltime}) ) is the fiducial mass accretion rate defined in Section 4.1."
+ Thus. fo is of order the accretion time onto the initial mass.," Thus, $t_0$ is of order the accretion time onto the initial mass."
+ Theupper pauel of Figure 7. plots (7) for and 0.5. obtained by nunerical iutegration of equation (71)).," Theupper panel of Figure \ref{fig:massvel} plots $\beta (\tau)$ for and 0.8, obtained by numerical integration of equation \ref{eqn:velocity}) )."
+ Also shown. in the lower panel. is the erowth of the nondimensional quantity M. the mass of the eravitating body relative to its initial value.," Also shown, in the lower panel, is the growth of the nondimensional quantity $M$, the mass of the gravitating body relative to its initial value."
+ As expected. the body slows down appreciably within au accretion time.," As expected, the body slows down appreciably within an accretion time."
+ This study has plvoted ou the close relatiouship between the dyuamical [riction force. Le.. the rausfer of linear momenttun [rom gas to a gravitating object. and the transfer of mass to that same object.," This study has pivoted on the close relationship between the dynamical friction force, i.e., the transfer of linear momentum from gas to a gravitating object, and the transfer of mass to that same object."
+ This relatiouship is emboclied in our central result. equation (58)).," This relationship is embodied in our central result, equation \ref{eqn:mdotv}) )."
+ From this equation. iu urn. we derived ai analytic expression for the force itself. equation (66)).," From this equation, in turn, we derived an analytic expression for the force itself, equation \ref{eqn:bondifric}) )."
+ We are now in a position to address a basic question raised in Section 2.1., We are now in a position to address a basic question raised in Section 2.1.
+ How are we justified in assuniug steady-state flow. when the gravitatiug body is coutinually deceleratiug?," How are we justified in assuming steady-state flow, when the gravitating body is continually decelerating?"
+ The auswer is that quasi-steady [low is established within a radius regi; over which the souud crossing time (resi/ 64) equals the time for the objects momentum to decrease appreciably (ALV/F)., The answer is that quasi-steady flow is established within a radius $r_{\rm crit}$ over which the sound crossing time $r_{\rm crit}/\cs$ ) equals the time for the object's momentum to decrease appreciably $MV/F$ ).
+ Recalling as ⋅ ∣⋅↽≻⋅↽≻ ∐⋜↕↕∌↥⊳∖∐∩⋅⋯⋜↕∐∠↩≼⊓∩⊋∣⊤↗∣∣≬⋡−∣⊲−⋜↕∐⇂⋃⊳∖⊔∩≺↵≺↽⋃⋜↕⋃∩∐↸⋅≻↖∖⊔⋅∖∖↽≺↵∐⋜↕∖⊽≺↵⋅∐∩∐≺⇂∐∐≺↵∐⊳∖↥≺≻∐⋜↕∐∖⊽⋅ ⋅d ⋅ ∣↡⋅ ⋅ whereporz).," Recalling that $F$ is normalized to $2\pi\,\rho_0\,c^2_{\rm s}\,r^2_{\rm s}$ and using equation \ref{eqn:mdotv}) ), we have, nondimensionally, where."
+. The latter quantity was implicitly assumed to be large from the start. when we neglected the selferavity of the gas.," The latter quantity was implicitly assumed to be large from the start, when we neglected the self-gravity of the gas."
+ The nondinensional mass accretion rate hovers near uuity for the entire evolution (recall Fig. 1)).," The nondimensional mass accretion rate hovers near unity for the entire evolution (recall Fig. \ref{fig:2ndg}) ),"
+ while AY itself starts at unity aud climbs., while $M$ itself starts at unity and climbs.
+ Heuce. the critical radius is much larger than rs. and our analysis is sell-cousistent.," Hence, the critical radius is much larger than $r_s$ , and our analysis is self-consistent."
+burst accretion episodes.,burst accretion episodes.
+" As already highlighted in BCGO09., episodic accretion can produce models of a given final mass that to various extents depart from standard non acereting models in terms of position in the HRD. Li abundance. and convective envelope mass."," As already highlighted in BCG09, episodic accretion can produce models of a given final mass that to various extents depart from standard non accreting models in terms of position in the HRD, Li abundance, and convective envelope mass."
+ Since the mechanisms responsible for the acereting disk formation and. possible outbursts depend on many environmental parameters (e.g. mass. temperature. angular momentum. magnetic field of the parent collapsing cloud. and thermal properties of the disk). itis certainly possible to produce various populations of stars at a similar age in young clusters and associations that depart from predictions of standard pre-main sequence models.," Since the mechanisms responsible for the accreting disk formation and possible outbursts depend on many environmental parameters (e.g. mass, temperature, angular momentum, magnetic field of the parent collapsing cloud, and thermal properties of the disk), it is certainly possible to produce various populations of stars at a similar age in young clusters and associations that depart from predictions of standard pre-main sequence models."
+ Therefore. our scenario does not predict a broad spread among. these stellar properties/signatures. because the level of departure entirely depends on the star or cluster properties anc aceretion histories.," Therefore, our scenario does not predict a broad spread among these stellar properties/signatures, because the level of departure entirely depends on the star or cluster properties and accretion histories."
+ Further observational effort is needed to explore these issues in more detail and to try to have better constraints on the processes of non steady accretior during the early stages of evolution., Further observational effort is needed to explore these issues in more detail and to try to have better constraints on the processes of non steady accretion during the early stages of evolution.
+ Among the multiple consequences of early episodic accretion. we emphasize the limited (to say the least) reliability of Li depletion as a age or cluster nembership indicator.," Among the multiple consequences of early episodic accretion, we emphasize the limited (to say the least) reliability of Li depletion as an age or cluster membership indicator."
+ We urge observers to re-analyze young cluster populations on the basis of the present results. since several genuine cluster members may have been eliminated because of their unexpectedly low Li abundance.," We urge observers to re-analyze young cluster populations on the basis of the present results, since several genuine cluster members may have been eliminated because of their unexpectedly low Li abundance."
+" Any information about the fraction of such Li depleted objects. with membership unambiguously assessed by radial velocity. could provide crucial insight into the occurrence and the intensity of non steady accretion,"," Any information about the fraction of such Li depleted objects, with membership unambiguously assessed by radial velocity, could provide crucial insight into the occurrence and the intensity of non steady accretion."
+ We also suggest that intense bursts of episodic accretion may have some link to planet formation.," We also suggest that intense bursts of episodic accretion may have some link to planet formation,"
+claborated by 2... were largely. in. place by the time Annie Cannon produced the Henry Draper Catalogue (?).. following earlier work by Fleming. Maury. and others (??).. in part under Pickering's supervision.,"elaborated by \citet{mkk}, were largely in place by the time Annie Cannon produced the Henry Draper Catalogue \citep[]{hdcat}, following earlier work by Fleming, Maury and others \citep{pick90b, maury97}, in part under Pickering's supervision."
+ In the A-type domain. Cannon's main classification criterion was the intensity of the A line (A33933.A)) with respect to L9 and the blend of He. | 44(A33968.A3).," In the A-type domain, Cannon's main classification criterion was the intensity of the $\;K$ line ) with respect to $\delta$ and the blend of $\epsilon$ + $\;H$ )."
+ ?. also identified 1e A line as à useful ciagnostic feature: it is the most conspicuous discriminant of temperature for A-type stars., \citet{m37} also identified the $\;$ $K$ line as a useful diagnostic feature; it is the most conspicuous discriminant of temperature for A-type stars.
+ We therefore investigated the use of the A line and the lle Lf blend as a temperature discriminant., We therefore investigated the use of the $K$ line and the $\epsilon$ $\;H$ blend as a temperature discriminant.
+ Although rere are. potentially. problems with the use of the A. line (discussed. in 4.3)). it has the advantage that a spectrum can be quickly and accurately. classified. even at ow signal-to-noise.," Although there are, potentially, problems with the use of the $K$ line (discussed in ), it has the advantage that a spectrum can be quickly and accurately classified, even at low signal-to-noise."
+ While other classification criteria exist. Lguch as the ratio of A44226 to Ale λΕΗΣΙ (7).. re weakness of most metal-line features in the spectra of (metal-deficient) SAIC stars makes them less useful for our »urposes. especially in view of the modest στα in the majority of our 2dLE spectra.," While other classification criteria exist, such as the ratio of $\;$ 4226 to $\;$ 4481 \citep{yam77}, the weakness of most metal-line features in the spectra of (metal-deficient) SMC stars makes them less useful for our purposes, especially in view of the modest s:n in the majority of our 2dF spectra."
+ The ratios of measured. Πες|££) line depths in. the Galactic stars are given in Table 1.. and are shown in Figure 4. as a function. of spectral as. expected. there is a well-defined. sequence.," The ratios of measured $K/{\rm H}\epsilon(+H)$ line depths in the Galactic stars are given in Table \ref{atargets}, and are shown in Figure \ref{ca} as a function of spectral as expected, there is a well-defined sequence."
+ However. there 1s. some overlap (e.g.. between Al and A2).," However, there is some overlap (e.g., between A1 and A2)."
+" In view of this. and the relatively poor sin of the SAIC data. we adopt a slightly less line classification unit. using only the ""dagger. subtypes for which standards are acdvaneec by 2: the results are sunmimarised in Table 3.. which alfords an almost completely ""clean segregation of the Galactic standards (Fig. 4))."," In view of this, and the relatively poor s:n of the SMC data, we adopt a slightly less fine classification unit, using only the `dagger' subtypes for which standards are advanced by \citet{mk73}; the results are summarised in Table \ref{aclass}, which affords an almost completely `clean' segregation of the Galactic standards (Fig. \ref{ca}) )."
+ The criteria in Table 3. have been adopted for classifying the SAIC stars., The criteria in Table \ref{aclass} have been adopted for classifying the SMC stars.
+ Even for quite poor-quality data. the AE ratio can be measured to better than ~0.05. which is generally sullicient. for classifications uncertain to better than one subtype at our adopted classification resolution.," Even for quite poor-quality data, the $K/{\rm
+H}\epsilon$ ratio can be measured to better than $\sim$ 0.05, which is generally sufficient for classifications uncertain to better than one subtype at our adopted classification resolution."
+ The scheme in Table 3. is largely in accord. with the original criteria of Cannon Pickering., The scheme in Table \ref{aclass} is largely in accord with the original criteria of Cannon Pickering.
+ Phe sensitivity of AVfM to spectral type falls olf towards cooler tempcratures. ancl for both AQ and FO types the A line is very strong GINAe=1): these subtypes are essentially inclistinguishable in the 2dE data. and we grouped them under the EO. label. (," The sensitivity of $K/{\rm
+H}\epsilon$ to spectral type falls off towards cooler temperatures, and for both A9 and F0 types the $K$ line is very strong $(K/{\rm
+H}\epsilon \simeq 1)$; these subtypes are essentially indistinguishable in the 2dF data, and we grouped them under the `F0' label. ("
+Still later spectral types are discussed in section 4.4)) The distinction between ADS and AY is relatively subtle. but is achieved consistently.,"Still later spectral types are discussed in section \ref{ftypes}) ) The distinction between A5 and A7 is relatively subtle, but is achieved consistently."
+ This is illustrated by the signal-weighted: mean spectra at cach spectral subtype shown in bil 3.., This is illustrated by the signal-weighted mean spectra at each spectral subtype shown in Fig. \ref{atypes3}.
+ Phe smooth progression of weak metal features with spectral type (e.g. Fe lzmd4173-78) demonstrates that the classifications are ab least internally consistent. while ακομαιοί. at A3/A5/AT is confirmed by the reversal of the ALI ratio between these subtypes.," The smooth progression of weak metal features with spectral type (e.g. $\;$ 4173-78) demonstrates that the classifications are at least internally consistent, while discrimination at A3/A5/A7 is confirmed by the reversal of the $K/{\rm H}\delta$ ratio between these subtypes."
+ Although relatively quick and simple to use. the Ca A (and I) lines are not without concerns for classification purposes.," Although relatively quick and simple to use, the Ca $K$ (and $H$ ) lines are not without concerns for classification purposes."
+ ? noted that using the calcium/hwdrogen ratio may introduce cllects into. the spectral-twpe classifications because the Balmer lines are sensitive to. luminosity., \citet{gg89b} noted that using the calcium/hydrogen ratio may introduce effects into the spectral-type classifications because the Balmer lines are sensitive to luminosity.
+ Llowever. the measurements in Table 1 show no obvious trend with luminosity class. as is illustrated by in Fie. 4..," However, the measurements in Table \ref{atargets} show no obvious trend with luminosity class, as is illustrated by in Fig. \ref{ca}."
+ Although the sampling is insullicienthy dense to rule out small differences between classes on average. the dispersion aoa given spectral subtype is. evidently much larger than such differences.," Although the sampling is insufficiently dense to rule out small differences between classes on average, the dispersion at a given spectral subtype is evidently much larger than such differences."
+ In any case. our 2dE sample. is overwhelmingly dominated by high-Iuminosity objects.," In any case, our 2dF sample is overwhelmingly dominated by high-luminosity objects."
+ lt has been known for a century that II&Ixeabsorplionlinesariscintheinterstellarmoediumt?)., It has been known for a century that $\;$ $H$ $K$ absorption lines arise in the interstellar medium \citep[][]{har04}.
+ E heinterstellarline. dispersionspectra.andceouldthereforeintroduceabiasintolhespectralelassif. ," The interstellar lines are not distinguishable in intermediate-dispersion spectra, and could therefore introduce a bias into the spectral classifications. ("
+0.1. the velocity dispersion. in. the line-of-sight interstellar material enhances the ellect of resonance-line absorption.),"Although dust extinction towards SMC stars is generally rather small, with $E(B-V) \la 0.1$ , the velocity dispersion in the line-of-sight interstellar material enhances the effect of resonance-line absorption.)"
+ To assess the likely impact of any contamination. reddening values were calculated as a crude surrogate for interstellar-line equivalent widths of Galactic targets (Table 1)).," To assess the likely impact of any contamination, reddening values were calculated as a crude surrogate for interstellar-line equivalent widths of Galactic targets (Table \ref{atargets}) )."
+ Reassurinely. there is no evidence of segregation between high- and low-reddening targets Fig. 5)).," Reassuringly, there is no evidence of segregation between high- and low-reddening targets Fig. \ref{calowe}) )."
+ ‘This is consistent with our expectations that (i) interstellar absorption is a relatively small contributor to the total Lf.fy equivalent widths for A-type stars. and (ii) interstellar Lf ," This is consistent with our expectations that (i) interstellar absorption is a relatively small contributor to the total $H, K$ equivalent widths for A-type stars, and (ii) interstellar $H$ "
+a chosen fiek of study.,a chosen field of study.
+ In the few cases where a studeut was a doible major. upper division GPA was calculated twice for each major aud recorded as separate data »oiuts.," In the few cases where a student was a double major, upper division GPA was calculated twice for each major and recorded as separate data points."
+ The set of upper division courses for ΛMathematics aud CIS (Computer and Information Sciences) were restricted to a subset of courses that are consklered rigorous aud are typically prereqtisites for post-g'aduate study., The set of upper division courses for Mathematics and CIS (Computer and Information Sciences) were restricted to a subset of courses that are considered rigorous and are typically prerequisites for post-graduate study.
+ As [9]iliiecl in greaer detail in the appendix. the ability to perform the analysis in this paper. »articularM7 he separation of GPA for upper division courses. required a combination of knowledge ofstudent records. acc‘ess {ο lists of majors over various vears. uetwork tools to extract tlie necessary data. software tools t«» interpret the data. aud an uucderstancing of data reduction of large datasets.," As outlined in greater detail in the appendix, the ability to perform the analysis in this paper, particularly the separation of GPA for upper division courses, required a combination of knowledge of student records, access to lists of majors over various years, network tools to extract the necessary data, software tools to interpret the data, and an understanding of data reduction of large datasets."
+ los oftje analysis involved in this study could not iive been produced by simple SQL commuauds O Veuilots databases suyported by student record software., Most of the analysis involved in this study could not have been produced by simple SQL commands to various databases supported by student record software.
+ In particular. typical access to student 'ecords in the current UO network environment is through webpages.," In particular, typical access to student records in the current UO network environment is through webpages."
+ Analysis of large numbers of stucdeut 'ecords requires the parsing aud interpretation of HTML. generating structured data from ustructured sources.," Analysis of large numbers of student records requires the parsing and interpretation of HTML, generating structured data from unstructured sources."
+ From the final sample. L179 (95% 4) had SAT scores.," From the final sample, 4,179 $95\%$ ) had SAT scores."
+ The rest bad ACT scores. which we converted to SAT values using standard ACT to SAT concordance tables.," The rest had ACT scores, which we converted to SAT values using standard ACT to SAT concordance tables."
+ In aggregate. our sample contains 8.085 students with SAT scores or converted ACT. scores.," In aggregate, our sample contains 8,085 students with SAT scores or converted ACT scores."
+ Of these students. 2.186 took he Math αιd Beadiug compoueus twice. 562 took them three times. 111 took them four times. 3l took thet Lover 5 times.," Of these students, 2,486 took the Math and Reading components twice, 562 took them three times, 114 took them four times, 31 took them over 5 times."
+ This allows us to analyze the ellect of multiple re-tests of the SAT ooking at tLe eall score versus he peak score aud the minium score versus the peak score., This allows us to analyze the effect of multiple re-tests of the SAT by looking at the mean score versus the peak score and the minimum score versus the peak score.
+ ΤΙ analysis DESlowh in Figures 1 ard 2. and demonstrates that. on average. re-taking the SAT ot liuproves sccwes by 20 to 30+) poitts per section.," This analysis is shown in Figures 1 and 2, and demonstrates that, on average, re-taking the SAT only improves scores by 20 to 30 points per section."
+ The use of average vs. peak SAT score has lit Impact ou ow subsectent analvsiS., The use of average vs. peak SAT score has little impact on our subsequent analysis.
+ Iu addition. from our initial search samples. we have 710 students with SAT scores a 'ecent. GRE scores (ypically graduating seniors applying [or adimissiou to graduate prograius).," In addition, from our initial search samples, we have 740 students with SAT scores and recent GRE scores (typically graduating seniors applying for admission to graduate programs)."
+ The composite SAT scores range from 800 to 1550 The SAT and GRE scores are highly correlated (with Ro valetes of rotelly 0.75 for both Math aud. Reading) aud a correlation slope near unity., The composite SAT scores range from 800 to 1550 The SAT and GRE scores are highly correlated (with R values of roughly 0.75 for both Math and Reading) and a correlation slope near unity.
+ This correlatlon Is sonewhat lower than the .86 obtained in Aneoll (1990). however that study was conducted before the 1905 re-centerine of the SAT.," This correlation is somewhat lower than the .86 obtained in Angoff (1990), however that study was conducted before the 1995 re-centering of the SAT."
+ Figure 3 displays the composite SAT scores (math plus critical reading) for students graduatiug from UO within oue of the twelve majors (1.120 students) versus those students who did. uot eraduate (2.[7E studeuts).," Figure 3 displays the composite SAT scores (math plus critical reading) for students graduating from UO within one of the twelve majors (4,420 students) versus those students who did not graduate (2,474 students)."
+ Note that this includes the eutire population of students who may have trauslerrecd to other schools with better (or different) programs. as well as students who left college," Note that this includes the entire population of students who may have transferred to other schools with better (or different) programs, as well as students who left college"
+Whether the metals found have a local origin associated with the Ly-a lines observed. or are due to an early episode of star formation (Population Hl) is stall unclear.,"Whether the metals found have a local origin associated with the $\alpha$ lines observed, or are due to an early episode of star formation (Population III) is still unclear."
+ The answer to this question mieht come from metal abundance studies in very low density clouds (CN(L Dx10H 7)., The answer to this question might come from metal abundance studies in very low density clouds (N(H $\leq 10^{14}$ $^{-2}$ ).
+ However. different investigations have reached contlicting conclusions.," However, different investigations have reached conflicting conclusions."
+ al.(1998) concluded that there is à sudden decrease of the C abundance at low column densities (ος—3.5 for NUL D<108 7). Cowie&Songaila(1998).," \citet{lu98} concluded that there is a sudden decrease of the C abundance at low column densities $\leq -3.5$ for N(H $< 10^{14}$ $^{-2}$ ). \citet{cow98},"
+. however. Claimed a roughly constant € IV abundance distribution all (he way down to Ly—a optical depths τν—a)~0.51.," however, claimed a roughly constant C IV abundance distribution all the way down to $-\alpha$ optical depths $\tau(\rm{Ly}-\alpha)\sim 0.5-1$."
+ More recently. Ellisonetal.(2000). have shown that the differences in redshift and. velocity. dispersion between Ly-a aud C IV. absorptions prevent us. al present. [rom answering of whether there is a unilorm degree of metal enrichment down to the lowest values of NOI I.," More recently, \citet{ell00} have shown that the differences in redshift and velocity dispersion between $\alpha$ and C IV absorptions prevent us, at present, from answering of whether there is a uniform degree of metal enrichment down to the lowest values of N(H I)."
+ An indirect wav to infer the possible existence of Population III and the nature of ils stellar mass function is to search for the imprints that could have been left in the IGM and/or in the subsequent generation of stars lormed from the material ejected. that is: (hechemical imprint.," An indirect way to infer the possible existence of Population III and the nature of its stellar mass function is to search for the imprints that could have been left in the IGM and/or in the subsequent generation of stars formed from the material ejected, that is: the imprint."
+ In (his paper we reexamine the issue of (he chemical pollution that could have originated Population IL., In this paper we reexamine the issue of the chemical pollution that could have originated Population III.
+ To do so. we make use of strictly zero metal (Z —0) stellar models in the mass range 3 m/AM.<200 (see below).," To do so, we make use of strictly zero metal $Z=0$ ) stellar models in the mass range $3\leq m/M_\odot\leq 200$ (see below)."
+ As an output. these models predict stellar vields for a number of significant isotopes which are the result of the nucleosvuthetic processes occurring during the evolution of the stars.," As an output, these models predict stellar yields for a number of significant isotopes which are the result of the nucleosynthetic processes occurring during the evolution of the stars."
+ These processes include those occurring in the thermal-pulsing AGB phase (TP-AGB) for the intermediate stellar mass range (3Sm/M. 8). in (ype II supernova (SNII) explosions (mr28 M.) and in very massive objects (VMO. m100 M. ).," These processes include those occurring in the thermal-pulsing AGB phase (TP-AGB) for the intermediate stellar mass range $3\la m/M_\odot\la 8$ ), in type II supernova (SNII) explosions $m\ga 8$ $_\odot$ ) and in very massive objects (VMO, $m\ga 100$ $_\odot$ )."
+ The paper is structured as follows: The observational constraints used to limit our model predictions are discussed in section 2., The paper is structured as follows: The observational constraints used to limit our model predictions are discussed in section 2.
+ The stellar vields for zero-metal stars are presented in section 3., The stellar yields for zero-metal stars are presented in section 3.
+ Section 4 describes the IME proposed for Population II., Section 4 describes the IMF proposed for Population III.
+ In section 5.," In section 5,"
+Throueh a quasi-linear force [ree method (Waugetal.2001).. Zhaoetal.(2008). sudied the uagnetlc topology sseletou of this active 'eglon.,"Through a quasi-linear force free method \citep{wang01}, , \citet{zhao08} studied the magnetic topology skeleton of this active region."
+ They foun the rising of the magnetic uull anc its spine. which is cousisteuo with the risiug o the extrapoatec [hx rope and the stroug twist ol the flux rope as shown ere.," They found the rising of the magnetic null and its spine, which is consistent with the rising of the extrapolated flux rope and the strong twist of the flux rope as shown here."
+ This strongly twisted flux roye provides a source for the shear of he post-Il:ure loop system at enhances the noi-poteutiality of associated iuagnetie coufiguratLOU., This strongly twisted flux rope provides a source for the shear of the post-flare loop system that enhances the non-potentiality of associated magnetic configuration.
+ Jingetal.(20009) st«lied tle temporal evelution of the 1iaglelic [ree energy for thisac ]ve 'eglol rou 21:00 to 23:00 UT o1 January 15 ard foixd that j| teids to iiCrease prior o tle ouset of lare 2., \citet{jing09} studied the temporal evolution of the magnetic free energy for this active region from 21:00 to 23:00 UT on January 15 and found that it tends to increase prior to the onset of flare 2.
+ The accuinulation o ‘the free euergy was mainly «'aused by the emergence 0o. Inagletic [lux and the roation of the stuspots. which continally traislers energy. below the plotospl‘e into he corona (Ixurolkawaetal.2002:Chene2010b).," The accumulation of the free energy was mainly caused by the emergence of magnetic flux and the rotation of the sunspots, which continually transfers energy below the photosphere into the corona \citep{kurokawa02,cheng10b}."
+. Tierefore. the post-flare corfigurat1 thay still accumulate magnetic ree energy in a short ime tha is respousible for the secprick erIption.," Therefore, the post-flare configuration may still accumulate magnetic free energy in a short time that is responsible for the second eruption."
+ li is worthy noting that the ionuologous fla'es can take place in the same active region with slinilar mo‘phologies and yattern of developmen (e.c..Choe&Cheng2000).. auc the 5ICcessive flares usualy occur in the uagneti¢ loops that are closely related bu uot in exactly the saue loops (e.g..Liuetal.2009).," It is worthy noting that the homologous flares can take place in the same active region with similar morphologies and pattern of development \citep[e.g.,][]{choe00}, and the successive flares usually occur in the magnetic loops that are closely related but not in exactly the same loops \citep[e.g.,][]{liu09}."
+. Whereas. he most significant difference. of he events iu this stucies [rom the genera homologous aud successive flares. is that the second CME/flare came directly rom the post-[lare oops of the first eruption but not the loops structures in the similar or other related locationis.," Whereas, the most significant difference, of the events in this studies from the general homologous and successive flares, is that the second CME/flare came directly from the post-flare loops of the first eruption but not the loops structures in the similar or other related locations."
+ I—1 conclusion. it is commonly viewed that a post-flare loop system evolves relatively ureventful alter tleir formation following a CME/Llare e‘uption.," In conclusion, it is commonly viewed that a post-flare loop system evolves relatively uneventful after their formation following a CME/flare eruption."
+ Except for some dynamic flows aloig the loops (e.g..Brekkeοἱal.1997:Qiuetal.1999).. the magnetic structure of the lox)ps is closer 1o poteutial than that before the eruption. hus unable to re-erupt in a short time afer a major eruption.," Except for some dynamic flows along the post-flare loops \citep[e.g.,][]{1997SoPh..175..511B,1999SoPh..190..153Q}, the magnetic structure of the loops is closer to potential than that before the eruption, thus unable to re-erupt in a short time after a major eruption."
+ The event studied bere provides ai exceptiona case aud may break such a point of view., The event studied here provides an exceptional case and may break such a point of view.
+ ]t shows that a post-Ilare magnetic coulieU‘ation can evolve quickly to a state tliat is favorable for another major CME/[Ilare eruption. as loig as there is Continual emergence of flux rope and/or rotation of sunspots.," It shows that a post-flare magnetic configuration can evolve quickly to a state that is favorable for another major CME/flare eruption, as long as there is continual emergence of flux rope and/or rotation of sunspots."
+ Post-flare loops cat re-[lare in a short time., Post-flare loops can re-flare in a short time.
+ Models ancl predictjon of solar flares (CMEs) should take into account sich an cobservational Lact., Models and prediction of solar flares (CMEs) should take into account such an observational fact.
+ We thaik the anonymous referee fo: his/her valuaje comments that helped to improve the paper., We thank the anonymous referee for his/her valuable comments that helped to improve the paper.
+ X.C1 NLD.D.. and Y.G. were supported by NSFC under grauts 10673001. 108252306. aud 1093230023 :ià NIXBRSFE. under erant. 2006CB806302.," X.C., M.D.D., and Y.G. were supported by NSFC under grants 10673004, 10828306, and 10933003 and NKBRSF under grant 2006CB806302."
+ .LZ. was supported by NSF erant ae NASA erant. NNCOSCOLIG. JJ. was supported by NSF under erant ATM aud ATM 07-16950., J.Z. was supported by NSF grant ATM-0748003 and NASA grant NNG05GG19G. J.J. was supported by NSF under grant ATM 09-36665 and ATM 07-16950.
+ T.W. was supported by DLR-erant 50 OC 0501.SOHO is a project of internatioial cooperation between ESA and NASA., T.W. was supported by DLR-grant 50 OC 0501.SOHO is a project of international cooperation between ESA and NASA.
+"models resemble the 'twisted torus' of Ciolfietal. (2009),, Lander&Jones (2009),, and Ciolfi,Ferrari,&Gualtieri(2010),, in that they consist of a potentially strong internal toroidal field, which is not observable directly and is confined inside the star, and an external field, which can be measured.","models resemble the `twisted torus' of \citet{cetal09}, , \citet{lj09}, , and \citet{cfg10}, in that they consist of a potentially strong internal toroidal field, which is not observable directly and is confined inside the star, and an external field, which can be measured."
+ The main difference is that we assume the star is stably stratified but non-barotropic (Reisenegger2009)., The main difference is that we assume the star is stably stratified but non-barotropic \citep{r09}.
+". Because of this, our models allow the poloidal and toroidal components to be adjusted independently of each other; in barotropic models (such as the twisted torus model), the poloidal and toroidal fields must obey certain relations to ensure that pressure is always a function of density purely [see also, e.g., Haskelletal. (2008)]]."," Because of this, our models allow the poloidal and toroidal components to be adjusted independently of each other; in barotropic models (such as the twisted torus model), the poloidal and toroidal fields must obey certain relations to ensure that pressure is always a function of density purely [see also, e.g., \citet{hetal08}] ]."
+" This means that our models can easily accommodate field configurations with 0.01<A0.8, the range found to be stable by the numerical simulations of Braithwaite&Nordlund(2006)."," This means that our models can easily accommodate field configurations with $0.01\lesssim\Lambda\lesssim 0.8$, the range found to be stable by the numerical simulations of \citet{bn06}."
+". As seen in Fig. 2,,"," As seen in Fig. \ref{eps},"
+ our calculated € is much lower than the LIGO and spin-down upper limits when the inferred dipole field strengths are input directly into Eq., our calculated $\epsilon$ is much lower than the LIGO and spin-down upper limits when the inferred dipole field strengths are input directly into Eq.
+ Note that the lower black dashed curve in Fig., Note that the lower black dashed curve in Fig.
+ 2 already assumes a strong toroidal field (A= 1073); a weaker toroidal field will generate a curve even lower down., \ref{eps} already assumes a strong toroidal field $\Lambda=10^{-3}$ ); a weaker toroidal field will generate a curve even lower down.
+" Next, we calculate how interior superconductivity (j,.= 10?) and accretion-induced screening or burial (red and blue bands in Fig. 2))"," Next, we calculate how interior superconductivity $\mu_r=10^{-3}$ ) and accretion-induced screening or burial (red and blue bands in Fig. \ref{eps}) )"
+ enhance c., enhance $\epsilon$.
+" This latter possibility is further explored in Fig. 4,,"," This latter possibility is further explored in Fig. \ref{xip},"
+" which shows how, in principle, a gravitational wave detection from a recycled pulsar with a certain surface field (measured from spin down) can be used to infer both A and shielding factor €."," which shows how, in principle, a gravitational wave detection from a recycled pulsar with a certain surface field (measured from spin down) can be used to infer both $\Lambda$ and shielding factor $\xi$."
+ Fig., Fig.
+" 4 also tells us the sets of parameters that are ruled out by current observations: for a given pulsar with a measured spin-down field strength, non-detection means that £ and A lie above that particular curve in Fig. 4.."," \ref{xip} also tells us the sets of parameters that are ruled out by current observations: for a given pulsar with a measured spin-down field strength, non-detection means that $\xi$ and $\Lambda$ lie above that particular curve in Fig. \ref{xip}."
+" For example, the upper solid curve of Fig. 2,,"," For example, the upper solid curve of Fig. \ref{eps},"
+" corresponding to £=1074 and A=107°, lies at the corner of the theoretically allowed shaded region in Fig."," corresponding to $\xi=10^{-4}$ and $\Lambda=10^{-3}$, lies at the corner of the theoretically allowed shaded region in Fig."
+ 4 and is therefore ruled out even for a non-superconducting star with a relatively weak observed field strength of 10* T (because it lies below the dashed blue curve in Fig. 4))., \ref{xip} and is therefore ruled out even for a non-superconducting star with a relatively weak observed field strength of $10^4$ T (because it lies below the dashed blue curve in Fig. \ref{xip}) ).
+ Is it possible that the magnetic field of a recycled pulsar is much stronger just below the surface than the inferred dipole field?, Is it possible that the magnetic field of a recycled pulsar is much stronger just below the surface than the inferred dipole field?
+" Krolik(1991) and Arons(1993),, motivated by the discrepancy between field strengths inferred from dipole spin down and from cyclotron line measurements of accreting X-ray pulsars, raised such a possibility."," \citet{k91} and \citet{a93}, motivated by the discrepancy between field strengths inferred from dipole spin down and from cyclotron line measurements of accreting X-ray pulsars, raised such a possibility."
+ They proposed that higher-order multipoles may exist close to the surface., They proposed that higher-order multipoles may exist close to the surface.
+" Arons(1993) then showed that surface field strengths ~10 times higher than the observed dipole field can account for the aforementioned discrepancy, as well as the location of the millisecond pulsars on the P-P diagram."," \citet{a93} then showed that surface field strengths $\sim 10^4$ times higher than the observed dipole field can account for the aforementioned discrepancy, as well as the location of the millisecond pulsars on the $P$ $\dot{P}$ diagram."
+" The surface field can also be masked by magnetic field burial, whereby accreted matter compresses the polar magnetic flux into a narrow belt around the equator."," The surface field can also be masked by magnetic field burial, whereby accreted matter compresses the polar magnetic flux into a narrow belt around the equator."
+ It has been shown thataccretionof >107?Mc is enough to alter, It has been shown thataccretionof $\gtrsim 10^{-5} M_\odot$ is enough to alter
+thus the contribution of massive stars to the ealactic chemical evolution is largo.,thus the contribution of massive stars to the galactic chemical evolution is large.
+ The black hole ]dnarv X-ray Nova S¢Ὁ (GRO J1655-10) may alsQ be affected. ]N a lypernova explosion. which can eject larec anount of Ti. S. aud Si aud leave a massive black hole as a coniaec ronuiaut.," The black hole binary X-ray Nova Sco (GRO J1655-40) may also be affected by a hypernova explosion, which can eject large amount of Ti, S, and Si and leave a massive black hole as a compact remnant."
+ Iu thils paper. we performed caculatioIs assululnes spherical svnuuetry and thus did not take account of nuon-sphDerical effects," In this paper, we performed calculations assuming spherical symmetry and thus did not take account of non-spherical effects."
+ For example. ifai jet Is associate wi idwpernovae. the explosive shock may be poiute« au stΤσο) iu this direction.," For example, if a jet is associated with hypernovae, the explosive shock may be pointed and stronger in this direction."
+" Iu this case, the lüeh exosion energies along the jet shoul have the properties of uucleosvuthesis in hvwper1ονας cliscussed above ar ortrogonal το the jet may be nuore like ordinary SUCTHOVAC."," In this case, the high explosion energies along the jet should have the properties of nucleosynthesis in hypernovae discussed above and orthogonal to the jet may be more like ordinary supernovae."
+ Iu such couditiois a hyovrnova could make ejecta which have a range of abuance ratios (as seen du Figure 6) in differcut directions.," In such conditions, a hypernova could make ejecta which have a range of abundance ratios (as seen in Figure 6) in different directions."
+" It those ejecta m various cürections interact with iutersellar uatter aud form mixed materials of various lass ra105, 8ars which could form from such interactions wouk have a range of metallicity and alttidanece ratios alo1g irou-peak elements."," If those ejecta in various directions interact with interstellar matter and form mixed materials of various mass ratios, stars which could form from such interactions would have a range of metallicity and abundance ratios among iron-peak elements."
+ For example. the ejecta in the jet direcion has larger amount of complete Si-burniug prodi cts{(1.6.. larger |Co/Foe]) aud are mixed with larger amount of hydrogen (1.c.. producing stualler [Fe/TI|) because of higher velocities than in other directions.," For example, the ejecta in the jet direction has larger amount of complete Si-burning products (i.e., larger [Co/Fe]) and are mixed with larger amount of hydrogen (i.e., producing smaller [Fe/H]) because of higher velocities than in other directions."
+ Then such explosions might be responsible for the abundance trends in the metal-poor stars., Then such explosions might be responsible for the abundance trends in the metal-poor stars.
+ Multi-dimensional simulations are needed to investigate these possibilities., Multi-dimensional simulations are needed to investigate these possibilities.
+ We would like to thank Drs., We would like to thank Drs.
+ D. Nordstrom. F. Spite. T. Tsuru. € Tsvrachan. CG. Brown. Ph.," B. Nordströmm, F. Spite, T. Tsuru, G. Israelian, G. Brown, Ph."
+ Podsiadlowski. P. Alazzali. aud D. Sclinidt for stimulating discussion on the observational indication of hyperuova nuucleosvuthliesis.," Podsiadlowski, P. Mazzali, and B. Schmidt for stimulating discussion on the observational indication of hypernova nucleosynthesis."
+ This work has been supported im part by the Crant-nr-AÀid for Scientific Research (12610233. 12710122) aud COE research (07CI22002) of the Japanese Ministry of Education. Scieuce. Culture. and Sports. Swiss National Scicuce Foundation evant 2000-53798.98. U.S. National Acronauties and Space Administration erant NÀCO-8105 aud the Joimt Tustitute for Heavy Tou Research. which has as nienber institutions the Universitv of Tennessee. Vouderbilt University. and the Oak Ridee National Laboratory.," This work has been supported in part by the Grant-in-Aid for Scientific Research (12640233, 12740122) and COE research (07CE2002) of the Japanese Ministry of Education, Science, Culture, and Sports, Swiss National Science Foundation grant 2000-53798.98, U.S. National Aeronautics and Space Administration grant NAG5-8405 and the Joint Institute for Heavy Ion Research, which has as member institutions the University of Tennessee, Vanderbilt University, and the Oak Ridge National Laboratory."
+ ORNL is managed by UT-Battclle. LLC. for the U.S. Departiment of Encrey. under coutract DE-AC(5-," ORNL is managed by UT-Battelle, LLC, for the U.S. Department of Energy under contract DE-AC05-00OR22725."
+We can evaluate this sum of products by considering another Fourier transform from p(£) back to w(s): Comparing this expression for the characteristic function to the original from Eq. (15)).,"We can evaluate this sum of products by considering another Fourier transform from $p(\xi)$ back to $\psi(s)$: Comparing this expression for the characteristic function to the original from Eq. \ref{eq:derivation_pxi_charfct}) ),"
+ we get Evaluation at s=0 yields so that together with Eq. (34)), we get Evaluation at $s=0$ yields so that together with Eq. \ref{eq:pxi_moments_norm}) )
+ we have shown that N=I., we have shown that $\mathcal{N}=1$.
+ To calculate the higher moments. we make use of the integral and of sum formulae of the type These follow from taking derivatives with respect to s» in Eq. (36))," To calculate the higher moments, we make use of the integral and of sum formulae of the type These follow from taking derivatives with respect to $s$ in Eq. \ref{eq:pxi_moments_charfcteq}) )"
+ and setting s=0., and setting $s=0$.
+ We have checked the results for mean. variance. skewness and kurtosis and have reproduced the results of the characteristic function approach. demonstrating the validity of the probability distribution function.," We have checked the results for mean, variance, skewness and kurtosis and have reproduced the results of the characteristic function approach, demonstrating the validity of the probability distribution function."
+ From the probability distribution function (22)). we can also directly calculate the cumulative distribution function. defined as F(é)=P(é&’>€).," From the probability distribution function \ref{eq:univar_derivation_pxifinal}) ), we can also directly calculate the cumulative distribution function, defined as $F(\xi) = P( \xi' > \xi)$."
+" For single poles only. it is given by Again. we use the notation 74,=—H(&H(C,) for the Heaviside factor. but this time note the extra term of -H(—£)."," For single poles only, it is given by Again, we use the notation $\mathcal{H}_n = H(\xi)H(C_n) - H(-\xi)H(-C_n)$ for the Heaviside factor, but this time note the extra term of $+ H(-\xi)$."
+ Analogous expressions in the presence of higher-order poles could be obtained by integrating the corresponding probability density (A3))., Analogous expressions in the presence of higher-order poles could be obtained by integrating the corresponding probability density \ref{eq:univar_derivation_pximultipole}) ).
+ In general. the probability distribution function we found ts a sum formula that needs to be evaluated numerically.," In general, the probability distribution function we found is a sum formula that needs to be evaluated numerically."
+ However. if the power spectrum of the underlying random field is a power law PUR)«[kl we can analytically find a more explicit expression for the univariate distribution function.," However, if the power spectrum of the underlying random field is a power law $P(k) \propto |k|^{-\nu}$, we can analytically find a more explicit expression for the univariate distribution function."
+" In the case of with A à normalisation constant. and for a separation of x2 0. we have C,=LA/(Acn) and the product factors are which 15 a special case of the infinite product family (2.p.754) The probability density function (at zero separation) is then where the field size L comes in through Eq. (4))."," In the case of with $A$ a normalisation constant, and for a separation of $x=0$ , we have $C_n = LA/\left(4\pi^2n^2\right)$ and the product factors are which is a special case of the infinite product family \citep[p. 754]{Prudnikov1986}
+ The probability density function (at zero separation) is then where the field size $L$ comes in through Eq. \ref{eq:derivation_sigma}) )."
+ We can now express the cumulative distribution function in terms of known functionsas, We can now express the cumulative distribution function in terms of known functionsas
+flux due to perigee passage of the compact object in a highly eccentric orbit.,flux due to perigee passage of the compact object in a highly eccentric orbit.
+ Finally. it should be noted that lies in the direction of the Sagittarius arm of our Galaxy. which ts a region rich of high mass/young stars. and therefore HMXBs.," Finally, it should be noted that lies in the direction of the Sagittarius arm of our Galaxy, which is a region rich of high mass/young stars, and therefore HMXBs."
+ This location could provide another indirect. support for being à HMXB. as proposed for 3 similar sources lying in the Norma arm (Revnivtsev 2003).," This location could provide another indirect support for being a HMXB, as proposed for 3 similar sources lying in the Norma arm (Revnivtsev 2003)."
+ This arm ts located about 2 kpe from the Sun. and if was associated with this region its luminosity as obtained from our spectral fit (10°-10°° ere s7! ) would be completely consistent with that of the aforementioned HMXB/neutron star binaries. as Vela Χ-Ι.," This arm is located about 2 kpc from the Sun, and if was associated with this region its luminosity as obtained from our spectral fit $^{35}$ $^{36}$ erg $^{-1}$ ) would be completely consistent with that of the aforementioned HMXB/neutron star binaries, as Vela X-1."
+ We have presented a detailed study of the hard X-ray properties of observed at different times with andRXTE., We have presented a detailed study of the hard X-ray properties of observed at different times with and.
+ From a_ well-sampled monitoring of the source in 2003 March-Mav. we deduced that spends most of its time in à low luminosity state. which likely corresponds to the state observed with on three occasions.," From a well-sampled monitoring of the source in 2003 March–May, we deduced that spends most of its time in a low luminosity state, which likely corresponds to the state observed with on three occasions."
+ The source spectrum is characteristic for thermal Comptonisation. and on one occasion we have evidence for a black body component in the spectrum.," The source spectrum is characteristic for thermal Comptonisation, and on one occasion we have evidence for a black body component in the spectrum."
+ From the comparison of the spectral properties of with those of other XRBs. we suggest that this system hosts a neutron star rather than a black The source spectra show evidence for a variable intrinsic absorption which indicate that the compact source is embedded in à dense cloud.," From the comparison of the spectral properties of with those of other XRBs, we suggest that this system hosts a neutron star rather than a black The source spectra show evidence for a variable intrinsic absorption which indicate that the compact source is embedded in a dense cloud."
+ This and the detection in all our spectra of a bright (and thin) iron line. whose flux is higher in the higher luminosity states points towards radial aecretion from a stellar wind.," This and the detection in all our spectra of a bright (and thin) iron line, whose flux is higher in the higher luminosity states points towards radial accretion from a stellar wind."
+ Therefore it is very likely that is a HMXB. with properties similar to those of other well known The arguments presented in the present study are. however. Cnly indicative. none of them being definite.," Therefore it is very likely that is a HMXB, with properties similar to those of other well known The arguments presented in the present study are, however, only indicative, none of them being definite."
+ In. particular the identification of counterparts at other wavelengths of the electromagnetic spectrum should allow one to truly confirm the nature of the system and/or the compact object., In particular the identification of counterparts at other wavelengths of the electromagnetic spectrum should allow one to truly confirm the nature of the system and/or the compact object.
+ Such à study is. however. not possible at the moment given the relatively large error on the position of the source.," Such a study is, however, not possible at the moment given the relatively large error on the position of the source."
+ Observations with high resolution X-ray satellites. such as orXMM-Newton. should permit a better position to be found. counterparts to be searched for. and possibly determine whether is indeed the same source as EXO 1912+097.," Observations with high resolution X-ray satellites, such as or, should permit a better position to be found, counterparts to be searched for, and possibly determine whether is indeed the same source as EXO 1912+097."
+ In addition such a study should permit one to obtain much better constraints on the absorption and line parameters., In addition such a study should permit one to obtain much better constraints on the absorption and line parameters.
+"important for every energy level) and take initial conditions along their unstable asymptotic curves, the integration of these asymptotic orbits can also reveal the morphological features of every energy level.","important for every energy level) and take initial conditions along their unstable asymptotic curves, the integration of these asymptotic orbits can also reveal the morphological features of every energy level."
+ In Fig., In Fig.
+" 20a, 20b we give an example of the projection on the configuration space of the apocentric and pericentric intersections (black points), respectively, of the 2D manifold of the —2:1 unstable periodic orbit corresponding to Ej;=—1135000<E;(£1) where the areas inside and outside corotation cannot communicate."," 20a, 20b we give an example of the projection on the configuration space of the apocentric and pericentric intersections (black points), respectively, of the 2D manifold of the $-2:1$ unstable periodic orbit corresponding to $E_j=-1135000<E_j(L_1)$ where the areas inside and outside corotation cannot communicate."
+" Note that the phase space in the area outside corotation, in the 2D approximation of the system, presents no tori or cantori with small holes, but only a chaotic sea with escapes to infinity."," Note that the phase space in the area outside corotation, in the 2D approximation of the system, presents no tori or cantori with small holes, but only a chaotic sea with escapes to infinity."
+ In each panel of Fig., In each panel of Fig.
+ 20 we have superimposed the density distribution (in color scale) of the real particles at the same Jacobi constant value., 20 we have superimposed the density distribution (in color scale) of the real particles at the same Jacobi constant value.
+" We observe that the apocentres do not coincide with the maxima of the density while the pericentres trace well the areas of the density maxima, but for slightly smaller radii."," We observe that the apocentres do not coincide with the maxima of the density while the pericentres trace well the areas of the density maxima, but for slightly smaller radii."
+ This is compatible to the information provided by Fig., This is compatible to the information provided by Fig.
+ 12 according to which for Jacobi constant values below Ej;(Li) and outside corotation the pericentres reveal the density maxima since they lie near low values of the plane velocity., 12 according to which for Jacobi constant values below $E_j(L_1)$ and outside corotation the pericentres reveal the density maxima since they lie near low values of the plane velocity.
+ Figure 21 is similar to Fig., Figure 21 is similar to Fig.
+" 20 but for a different Jacobi constant E;—-1120000, which is just above Ej(Li), and therefore the area inside corotation has just communicated with the area outside corotation."," 20 but for a different Jacobi constant $E_j$ =-1120000, which is just above $E_j(L_1)$, and therefore the area inside corotation has just communicated with the area outside corotation."
+" In this energy level and in the 2D approximation there exist unstable periodic orbits of the PLi, PL» and —1:1 families."," In this energy level and in the 2D approximation there exist unstable periodic orbits of the $PL_1$, $PL_2$ and $-1:1$ families."
+ In Fig., In Fig.
+" 21a,b we plot the apocentres and the pericentres, respectively, of all the asymptotic orbits (having initial conditions on the unstable asymptotic curves) of all the periodic orbits mentioned above, superimposed on the density distribution of the real particles in the specific Jacobi constant value."," 21a,b we plot the apocentres and the pericentres, respectively, of all the asymptotic orbits (having initial conditions on the unstable asymptotic curves) of all the periodic orbits mentioned above, superimposed on the density distribution of the real particles in the specific Jacobi constant value."
+" In this case we see that the pericentres together with some parts of the apocentres (mostly of the PLi;, PL» families) trace well the density maxima of the density distribution."," In this case we see that the pericentres together with some parts of the apocentres (mostly of the $PL_1$, $PL_2$ families) trace well the density maxima of the density distribution."
+ This is compatible to the information provided by Fig., This is compatible to the information provided by Fig.
+ 13 according to which the pericentres and the apocentres for Jacobi constant values above E;(Li) reveal the density maxima when they lie near low values of the plane velocity., 13 according to which the pericentres and the apocentres for Jacobi constant values above $E_j(L_1)$ reveal the density maxima when they lie near low values of the plane velocity.
+" The stickiness effect along the asymptotic manifolds of the unstable periodic orbits can last for a number of dynamical times reinforcing the spiral structure, as it has been described in the previous sections, but then the chaotic orbits can escape to very large distances and finally they may escape to infinity."," The stickiness effect along the asymptotic manifolds of the unstable periodic orbits can last for a number of dynamical times reinforcing the spiral structure, as it has been described in the previous sections, but then the chaotic orbits can escape to very large distances and finally they may escape to infinity."
+ This diffusion is relatively slow and its time scale can be compared to the age of the Universe., This diffusion is relatively slow and its time scale can be compared to the age of the Universe.
+ In this section we study the way these chaotic orbits escape to infinity., In this section we study the way these chaotic orbits escape to infinity.
+After more than a decade of continuous improvements in the accuracy of cosmological observations which has led to the establishment of a broadly accepted: representation of our Universe known as theAlodel (COAI) we are now entering the epoch of precision cosmology.,After more than a decade of continuous improvements in the accuracy of cosmological observations – which has led to the establishment of a broadly accepted representation of our Universe known as the (CCM) – we are now entering the epoch of precision cosmology.
+ The great. wealth of high-precision cosmological data expected throughout the next few vears ollers the exciting prospect of tightly constraining the parameters of the CCAL or. possibly detecting. deviations [rom its predictions., The great wealth of high-precision cosmological data expected throughout the next few years offers the exciting prospect of tightly constraining the parameters of the CCM or possibly detecting deviations from its predictions.
+ The unprecedented: resolution of the Planck satellite (Adeetal.20011). in measuring the angular temperature Iluetuations and the polarization of the Cosmic Microwave Dackground (CAIB) radiation. will extract a significant amount of new information bevond that vielded by previous CAIB experiments. alfording very tight. constraints on the initial conditions of the Universe.," The unprecedented resolution of the Planck satellite \citep[][]{Ade:2011ah} in measuring the angular temperature fluctuations and the polarization of the Cosmic Microwave Background (CMB) radiation will extract a significant amount of new information beyond that yielded by previous CMB experiments, affording very tight constraints on the initial conditions of the Universe."
+ At the other end of the cosmic expansion history. a low redshifts. present and future surveys measuring the clustering of sources (Teamctal.2011) and eravitationa lensing ellects (Masseyctal.2007:Teamet2011). wil greatly improve our knowledge of the spatial distribution of barvonic ancl cold. dark matter (CDM). in. the. loca Universe. and of its time evolution. providing new tests of 16 gravitational instability processes driving the growth of COsDmIIC slpuctures.," At the other end of the cosmic expansion history, at low redshifts, present and future surveys measuring the clustering of sources \citep[][]{Euclid-r} and gravitational lensing effects \citep[][]{Massey:2007gh,Euclid-r} will greatly improve our knowledge of the spatial distribution of baryonic and cold dark matter (CDM) in the local Universe, and of its time evolution, providing new tests of the gravitational instability processes driving the growth of cosmic structures."
+ In particular. one of the most mysterious phenomena characterizing the Iow-redshift. Universe is the appearance of a Dark Energy. (DIE) component capable of driving the observed. present acceleration of the cosmic expansion (ItiessAstieretal.2006:Ixowalskict 2008).. and: required to also explain a number of other observations. the angular power spectrum of CMD temperature anisotroples (IXomatsuetal.2009.2011).. the evolution of the number counts of massive galaxy clusters as a function of recdshift (Dorgani2006:Vikhlininetal.2009:Mantz2010).. the angular correlation of galaxies in large galaxy surveys (Percivaletal.2001:Cole2005:Reid2010).. or the observed. scale of the Barvon Acoustic Oscillations," In particular, one of the most mysterious phenomena characterizing the low-redshift Universe is the appearance of a Dark Energy (DE) component capable of driving the observed present acceleration of the cosmic expansion \citep[][]{Riess_etal_1998,Perlmutter_etal_1999,Schmidt_etal_1998,snls,kowalski_etal_2008}, and required to also explain a number of other observations, the angular power spectrum of CMB temperature anisotropies \citep[][]{wmap5,wmap7}, the evolution of the number counts of massive galaxy clusters as a function of redshift \citep[][]{Borgani_2006,Vikhlinin_etal_2009b,Mantz_etal_2010}, the angular correlation of galaxies in large galaxy surveys \citep[][]{Percival_etal_2001,Cole_etal_2005,Reid_etal_2010}, or the observed scale of the Baryon Acoustic Oscillations"
+spectra (indicated by the unbroken power-law continua out to 2100 keV. see e.g. MeClintock Remillard 2006: although see c.g. Zdziarski et al.,"spectra (indicated by the unbroken power-law continua out to $> 100$ keV, see e.g. McClintock Remillard 2006; although see e.g. Zdziarski et al."
+ 2005 for the ULX-like behaviour of GRS 1915|105)., 2005 for the ULX-like behaviour of GRS 1915+105).
+ We confirm this physical cilference by fixing the plasma temperature at 50 keV in our mocels. and measuring the resulting change in fit quality.," We confirm this physical difference by fixing the plasma temperature at 50 keV in our models, and measuring the resulting change in fit quality."
+ In all three cases the fit is poorer with a hot corona. by Ay? of 134. 42 and 15 respectively for one more degree of freedom.," In all three cases the fit is poorer with a hot corona, by $\Delta\chi^2$ of 134, 42 and 15 respectively for one more degree of freedom."
+ This shows that a cool corona provides a improvement according to the F-test for the high Ilux bin. and a substantially more significant improvement for the other bins. compelling evidence that the corona does not appear similar to that in standard sub-IEddington states.," This shows that a cool corona provides a improvement according to the F-test for the high flux bin, and a substantially more significant improvement for the other bins, compelling evidence that the corona does not appear similar to that in standard sub-Eddington states."
+ Additionally. the Compton tail in a sub-IEXddington state is highly variable on short. timescales providing an unambiguous test. for such a model.," Additionally, the Compton tail in a sub-Eddington state is highly variable on short timescales providing an unambiguous test for such a model."
+ We characterise the variability using the excess rms of the 3 10 keV lightcurve in each Dux bin (binned to 250 s). Le. the variability above the Poisson (white) noise level of the lightcurve (see Ecelson et al.," We characterise the variability using the excess rms of the 3 – 10 keV lightcurve in each flux bin (binned to 250 s), i.e. the variability above the Poisson (white) noise level of the lightcurve (see Edelson et al."
+ 2002)., 2002).
+ We find that there is constrained variability above 3 keV with 30 upper limits of «θα. «Ss and «114 for the three Dux. bins respectively.," We find that there is constrained variability above 3 keV with $\sigma$ upper limits of $<$, $<$ and $<$ for the three flux bins respectively."
+ ‘This strongly argues against a two-component sub-IEEddington model for the data., This strongly argues against a two-component sub-Eddington model for the data.
+ Indeed. we note that the combination of a hot disc (in 2/3 fits). and a cool. optically thick and invariant corona. demonstrates that this ULNA cannot be described. by the cool dise plus power-law continuum: tov mocel previously used to infer INIBIISs in ULXs.," Indeed, we note that the combination of a hot disc (in 2/3 fits), and a cool, optically thick and invariant corona, demonstrates that this ULX cannot be described by the cool disc plus power-law continuum toy model previously used to infer IMBHs in ULXs."
+ Ifthe data were to be described solely by disc emission then we would not expect there to. be any variability on anvthing other than the longest timescales (Wilkinson Uttley 2009)., If the data were to be described solely by disc emission then we would not expect there to be any variability on anything other than the longest timescales (Wilkinson Uttley 2009).
+ As we have substantial evidence against the Ovo-component sub-Ecddington model. and. are now interested. in the variability behaviour of the emission. as a whole. we extract the Fourier-frequencey dependent: power density spectra (PDS) over the full energy bandpass.," As we have substantial evidence against the two-component sub-Eddington model and are now interested in the variability behaviour of the emission as a whole, we extract the Fourier-frequency dependent power density spectra (PDS) over the full energy bandpass."
+ Weng et al. (, Weng et al. (
+2009) report the lack of high frequeney: variability in the inclivicual observations.,2009) report the lack of high frequency variability in the individual observations.
+ Here however. we can extend this to longer timescales by selecting the shortest segment of continuous lighteurve and taking integer number of intervals of this length across the remaining observations.," Here however, we can extend this to longer timescales by selecting the shortest segment of continuous lightcurve and taking integer number of intervals of this length across the remaining observations."
+ This can require certain finessing to obtain the maximum. amount of available data and. broadest frequency bandpass (as this goes from 1longest available individual segment length] o bin size])., This can require certain finessing to obtain the maximum amount of available data and broadest frequency bandpass (as this goes from 1/[longest available individual segment length] to $\times$ bin size]).
+ In Figure 1 we present the average »ower density spectrum (PDS - extracted using the POWSPEC) for each Bux bin. normalised to rms? units ollowing geometrical re-binning with white noise included.," In Figure 1 we present the average power density spectrum (PDS - extracted using the ) for each flux bin, normalised to $^2$ units following geometrical re-binning with white noise included."
+ Quite clearly there is no constrained: variability over any of the available frequency bins. consistent with the source jwing suppressed red noise as seen in a handful of ULAs (lloil et al.," Quite clearly there is no constrained variability over any of the available frequency bins, consistent with the source having suppressed red noise as seen in a handful of ULXs (Heil et al."
+ 2009). and the emission originating in the accretion disc.," 2009), and the emission originating in the accretion disc."
+ llowever. a thin disc is a poor description. of the spectral cata as the Wien. tail is not broad. enough. to fit the high energy emission. leading to poor [fits (X2 of 1441.5/658. 931.4/617 and 650.7/481 lor an absorbed multi-colour cise blackbody spectrum in fit to each Εαν bin respectively).," However, a thin disc is a poor description of the spectral data as the Wien tail is not broad enough to fit the high energy emission, leading to poor fits $\chi^2$ of 1441.5/658, 931.4/617 and 650.7/481 for an absorbed multi-colour disc blackbody spectrum – in – fit to each flux bin respectively)."
+ Disc emission that has been smeared by relativistic cllects is considerably broader ancl so we also fitted the spectral cata with the modelKERRBB., Disc emission that has been smeared by relativistic effects is considerably broader and so we also fitted the spectral data with the model.
+ The data and best fitting models are shown in Figure 2. and the fit. parameters are detailed in Table 2.," The data and best fitting models are shown in Figure 2, and the fit parameters are detailed in Table 2."
+ Even in this extremely relativistically broacdencd case we still find. the models. poorly. describe the data. with obvious excesses above the model at. high energies.," Even in this extremely relativistically broadened case we still find the models poorly describe the data, with obvious excesses above the model at high energies."
+ It would therefore appear that none of the sub-Iddington models can describe both the broad shape of the spectra and also the lack of high energy. variability., It would therefore appear that none of the sub-Eddington models can describe both the broad shape of the spectra and also the lack of high energy variability.
+ Examples of Ecclington accretion rates (i.e. accretion rates closely approximating those expected at the Ecddington limit [or an object) have been inferred. in the outburst of several Galactic N-rav binaries. for cxaniple ΜΗ (νο (Zvvcki. Done Smith 1909). V4641 Ser (Revnivisey et al.," Examples of Eddington accretion rates (i.e. accretion rates closely approximating those expected at the Eddington limit for an object) have been inferred in the outburst of several Galactic X-ray binaries, for example V404 Cyg (Żyycki, Done Smith 1999), V4641 Sgr (Revnivtsev et al."
+ 2002) and the neutron star svstem Cir N-1 (Done Cierlisski 2003)., 2002) and the neutron star system Cir X-1 (Done Gierlińsski 2003).
+ Llowever. ULXs are observed to be persistenthy Iuminous. which would suggest a closer analogy to a long-lived outburst system such as GRS 1915|105 (see Bemillaed AleClintock 2006). althoug= his source is also known o undergo dramatic state changes on short timescales (e.g. Belloni et al.," However, ULXs are observed to be persistently luminous, which would suggest a closer analogy to a long-lived outburst system such as GRS 1915+105 (see Remillard McClintock 2006), although this source is also known to undergo dramatic state changes on short timescales (e.g. Belloni et al."
+ 2000: Middleton et al., 2000; Middleton et al.
+ 2006)., 2006).
+ Taking a theoretical approach. we would expect that. as the Ecelington limit is neared. the properties of the low through the disc would change as advection becomes more important as a method for removing energv from he Low (Mineshige ct al.," Taking a theoretical approach, we would expect that, as the Eddington limit is neared, the properties of the flow through the disc would change as advection becomes more important as a method for removing energy from the flow (Mineshige et al."
+ 2000: Abramowicz ct al., 2000; Abramowicz et al.
+ LOSS)., 1988).
+ We would also expect that. where the How is highly illuminated. material would be driven olf the cise surface » radiation pressure. with the massive outllow forming a photosphere above the disc (see Poutanen et al.," We would also expect that, where the flow is highly illuminated, material would be driven off the disc surface by radiation pressure, with the massive outflow forming a photosphere above the disc (see Poutanen et al."
+ 2007)., 2007).
+ Figure G left panel shows 5(T.g4) and S(r/—Yep as €.;Ep Changes for Ου! ASA., Figure \ref{fig:signif_obs} left panel shows $S({T_{\rm eff\it A}})$ and $S({r-i})_{EB}$ as $\epsilon_{r-iEB}$ changes for $_{EB}$ +AS3.
+" As e,pep decreases iu precision. S(r—/)gg decreases from 1.0 to 0.6. while au increase m S(T) frou 0.6 to 0.8 is also ποσα. perhaps This is sinilar to the sinele star case but in the binary svstem Jug remains relatively iuore important because the fended (r/)gp provides less information (see Fie."," As $\epsilon_{r-iEB}$ decreases in precision, $S({r-i})_{EB}$ decreases from 1.0 to 0.6, while an increase in $S({T_{\rm eff\it A}})$ from 0.6 to 0.8 is also seen, perhaps This is similar to the single star case but in the binary system $_A$ remains relatively more important because the blended $(r-i)_{EB}$ provides less information (see Fig."
+ 3 oft diagram)., \ref{fig:signif_obs_s} left diagram).
+ The right halfof Fieure 6 shows SR4) as its precision is decreased. for OSgp and OSpp! ASL., The right half of Figure \ref{fig:signif_obs} shows $S({R_A})$ as its precision is decreased for $_{EB}$ and $_{EB}$ +AS1.
+ the same tvpe of information., the same type of information.
+ The observables that showed most sensitivity to a change in their precisions for different sets for cither a single star or a binary svstem are FR. μι the photometric colors and an iceutified mode.," The observables that showed most sensitivity to a change in their precisions for different sets for either a single star or a binary system are $R$, , the photometric colors and an identified mode."
+ Iu this section we study the uncertainties (Eq. 1)), In this section we study the uncertainties (Eq. \ref{eqn:uncertainties}) )
+ im the global paralucters of the (primary) pulsating star as these observables are improvedincluded for both a single star aud a binary svsteini., in the global parameters of the (primary) pulsating star as these observables are improved/included for both a single star and a binary system.
+" To successfullv use a single ideutified mode as a ""seine probe”. we need to know the fundamental stellar paralucters M. r. X. Z aud e of the pulsating star with eood precision."," To successfully use a single identified mode as a “seismic probe”, we need to know the fundamental stellar parameters $M$, $\tau$, $X$, $Z$ and $v$ of the pulsating star with good precision."
+ Here we diseuss only the uncertainties in the first four. because the determination of ce is incdependeut of the determination of the others (the covariance off-diagonal elements ~ 0) and independent of most of the observable errors. except / aud c4sin.," Here we discuss only the uncertainties in the first four, because the determination of $v$ is independent of the determination of the others (the covariance off-diagonal elements $\sim0$ ) and independent of most of the observable errors, except $i$ and $v_A\sin i$."
+ Figue 7 shows the theoretical uucertauties in AL (top). ((second). .X (third) and Z (lower panels) as we increase the precision iu A (left panels) aud both of the photometric colors (right paucls).," Figure \ref{fig:binsglpar} shows the theoretical uncertainties in $M$ (top), (second), $X$ (third) and $Z$ (lower panels) as we increase the precision in $R$ (left panels) and both of the photometric colors (right panels)."
+ The latter is denoted bve; ;aud for the binary it implies the bleuded systems colors., The latter is denoted by $\epsilon_{i-z}$ and for the binary it implies the blended system's colors.
+" The dashed aud solid lines represcut the single aud binary svstenm respectively, aud the observable set with the identified mode has dots overplotted on the lues."," The dashed and solid lines represent the single and binary system respectively, and the observable set with the identified mode has dots overplotted on the lines."
+ The results for OS (both OSs aud OSrgp) aud OS|ASL are represented ou the left paucls. aud those for OS|AS2 and OS|AS3 (incliding photometric iuforiiation) on the vieht.," The results for OS (both $_S$ and $_{EB}$ ) and OS+AS1 are represented on the left panels, and those for OS+AS2 and OS+AS3 (including photometric information) on the right."
+ Both the left aud right paucls show the same v-axcs scales for comparison., Both the left and right panels show the same y-axes scales for comparison.
+ It is very clear that the oscillation mode plays a very different role in the sinele star and the binary system., It is very clear that the oscillation mode plays a very different role in the single star and the binary system.
+ Excludius it from the single star constraiuts. without photometry. vields uucertaities > and independent of eg: the coustraimts are inadequate o determine the stellar model.," Excluding it from the single star constraints, without photometry, yields uncertainties $>$ and independent of $\epsilon_{R}$; the constraints are inadequate to determine the stellar model."
+ Tucliding the color constraiuts also vields large uncertainties (> 30%)) although not eutielv independent of e;.., Including the color constraints also yields large uncertainties $>$ ) although not entirely independent of $\epsilon_{i-z}$.
+ When the uode is included. the uncertainties miprove rapidlv as he errors iu both types of observables improve.," When the mode is included, the uncertainties improve rapidly as the errors in both types of observables improve."
+ In fact. J{AL) becomes strictly depeudent on ἐν as the error in he radius improves to «3 .," In fact, $\sigma(M)$ becomes strictly dependent on $\epsilon_R$ as the error in the radius improves to $<2$ –."
+ There is Little difference secu din the parameter uncertainties for the binary svstem when the oscillation node is inchided/exchided as a coustrait., There is little difference seen in the parameter uncertainties for the binary system when the oscillation mode is included/excluded as a constraint.
+ The binary system without auv seismic data provides simular or etter constraiuts on all of the parameters than the suele star an ideutified mode. except im some exceptional cases (see below).," The binary system without any seismic data provides similar or better constraints on all of the parameters than the single star an identified mode, except in some exceptional cases (see below)."
+ The fact that iucludiug a anode with the binary system coustraits does not cad to iunprovenieuts in the pulsating stars parameter uncertainties mudicates that the binary constraints alone nay provide sufficient coustraiuts on the model to he identified iode in a different way., The fact that including a mode with the binary system constraints does not lead to improvements in the pulsating star's parameter uncertainties indicates that the binary constraints alone may provide sufficient constraints on the model to the identified mode in a different way.
+ or both a(Z) aud o(7) for OS|ASS. In some cases the sinele star provides slightly better constraints than the binary.," We note that for both $\sigma(Z)$ and $\sigma(\tau)$ for OS+AS3, in some cases the single star provides slightly better constraints than the binary."
+ This happens ouly for very precisely measured (unubleuded) colors. while iucludiug an oscillation node as a constraint.," This happens only for very precisely measured (unblended) colors, while including an oscillation mode as a constraint."
+ Finally. oue should also note that cCX) 2 is not a good coustraint.," Finally, one should also note that $\sigma(X)$ $\gtrsim$ is not a good constraint."
+ The absolute value of the parameter is 0.700. and a 0.070 error on this value produces uo meaningful coustraiut on X.," The absolute value of the parameter is 0.700, and a 0.070 error on this value produces no meaningful constraint on $X$ ."
+ The inatrix from the SV decomposition describes the correlations among the parameters. while the singular values scale cach of the vectors V; to produce tlie u-dimensional error ellipse axes.," The matrix from the SV decomposition describes the correlations among the parameters, while the singular values scale each of the vectors $V_j$ to produce the n-dimensional error ellipse axes."
+ The black/erey ellipses correspond. to the sets of errors in the third cohuuu of Table 2. for OSgp |AS2/AS3., The black/grey ellipses correspond to the sets of errors in the third column of Table \ref{tab:obsers} for $_{EB}$ +AS2/AS3.
+ Tuchiding the oscillation mode does not reduce the error ellipse. and this reiuforces the possibility of using this extra Information to learn something else about the star.," Including the oscillation mode does not reduce the error ellipse, and this reinforces the possibility of using this extra information to learn something else about the star."
+ Using equation (2)) we calculate the theoretical nucertaitics iu the model effective temperature ancl model huuinositv., Using equation \ref{eqn:dscuti_er}) ) we calculate the theoretical uncertainties in the model effective temperature and model luminosity.
+ Figure 9 shows the approximate theoretical error boxes for these quantities for a single star (OSs |ASI. lines} aud a binary system (OSrp lines).," Figure \ref{fig2} shows the approximate theoretical error boxes for these quantities for a single star $_S$ +AS1, ) and a binary system $_{EB}$ )."
+ For the single staran identified mode is Πιοποσα becauseFig., For the single staran identified mode is included becauseFig.
+ 7 aud Sect., \ref{fig:binsglpar} and Sect.
+ ?? showed that the," \ref{sec:4.1}
+ showed that the"
+to measure tangential shear profile for cach subcluup candidate.,to measure tangential shear profile for each subclump candidate.
+" We therefore measure a profile of tangerial shear components euseiibling five subceliuup candidaes,", We therefore measure a profile of tangential shear components ensembling five subclump candidates.
+ A ass nieasureimnenut with stacked taugeutial distortion profile is complenmeΠαιν to the ο statistics ) (85)). because different &rear catalogue is used in two measurements.," A mass measurement with stacked tangential distortion profile is complementary to the $\zeta_c$ statistics \ref{sec:submass}) ), because different shear catalogue is used in two measurements."
+ Iu t1ο Qs statistics. source ealaxics outside a eiven radit sas used. while. iu the tangential shear iieasurement. source galaxies from inner to οιter radius is mdependenuIvy available.," In the $\zeta_c$ statistics, source galaxies outside a given radius is used, while, in the tangential shear measurement, source galaxies from inner to outer radius is independently available."
+ We here exclude two dark halos associated with cD ealaxics in order to avoid a contamination of lesing distortion caused by the main cluster., We here exclude two dark halos associated with cD galaxies in order to avoid a contamination of lensing distortion caused by the main cluster.
+ The candidate 5 at which the projection effect is sienificaut 1s also ignored., The candidate 5 at which the projection effect is significant is also ignored.
+" The ceuter for cach subcblup is chosen as the same position of the Mz, measurements (853).", The center for each subclump is chosen as the same position of the $M_{\zeta_c}$ measurements \ref{sec:submass}) ).
+" The averaged tangential shear distortions of source eaOs:axies. (gi(0,). is calculated in the circular auuulus of the sae radius. based on the same procedure as refsecinass.."," The averaged tangential shear distortions of source galaxies, $\langle g_+ \rangle
+(\theta_n)$, is calculated in the circular annulus of the same radius, based on the same procedure as \\ref{sec:mass}."
+" The vpieal projected distance between a ceiter of a stackec tanecutial shear profile aud the main chster center. (0)U, Jas obtained with a weielt function of lensing signals. (ar)2-Misi/2020Mig)17! |l. where (yylj is the leusine signal (83)) for cach SHοπήςand 6,jo an angular separation between"," The typical projected distance between a center of a stacked tangential shear profile and the main cluster center, $\langle \theta_{{\rm off}} \rangle$, is obtained with a weight function of lensing signals, $\langle \theta_{{\rm off}} \rangle^2=\sum_j \langle g_{+,j} \rangle^2 \theta_{{\rm off},j}^2/ \sum_j
+\langle g_{+,j} \rangle ^2\sim17\farcm4$ , where $\langle g_{+,j} \rangle$ is the lensing signal \ref{sec:mass}) ) for each subclumpand $\theta_{{\rm off},j}$ is an angular separation between"
+2.4 Hadron-nucleus collision,and the rate of the momentum loss of $h$ is set to be $P(z)=0.25z^{0.25-1}$ from the data on the
+Seyfert nucleus.,Seyfert nucleus.
+" The Chandra—measured position of the nuclear source is a.(2000)=LLM3™oge9[. 5.(2000)=—6572021""0."," The –measured position of the nuclear source is $\alpha_{\mathrm x}(2000) = 14^{\mathrm h} 13^{\mathrm m}
+09.\!^{\mathrm s}94$ , $\delta_{\rm x}(2000) = -65^{\circ} 20^{\prime}
+21.\!^{\prime\prime}0$."
+" This position is 176 away from the optical position OFLl. d,(2000)=—62520'21""|+ 0""7: Freeman οἱ al."," This position is $1.\!^{\prime\prime}6$ away from the optical position $\alpha_{\mathrm o}(2000) = 14^{\mathrm h}
+13^{\mathrm m} 09.\!^{\mathrm s}7 \pm 0.\!^{\mathrm s}1$ , $\delta_{\mathrm o}(2000) = -65^{\circ} 20^{\prime}
+21.\!^{\prime\prime}4 \pm 0.\!^{\prime\prime}7$ ; Freeman et al."
+" 1977) and 0/72 away [rom the radio position fa,{2000)=11513900795+0702. 5,(2000)=—65720'21""2+0""1: Creenbill 2001)."," 1977) and $0.\!^{\prime\prime}2$ away from the radio position $\alpha_{\mathrm r}(2000) = 14^{\mathrm h} 13^{\mathrm m}
+09.\!^{\mathrm s}95 \pm 0.\!^{\mathrm s}02$, $\delta_{\mathrm r}(2000)
+= -65^{\circ} 20^{\prime} 21.\!^{\prime\prime}2 \pm
+0.\!^{\prime\prime}1$; Greenhill 2001)."
+ The X-ray and radio positious are thus in agreemeut., The X-ray and radio positions are thus in agreement.
+ The 1nner 31725x (600x600 pc) region of the field is shown in more detail in Figure 2..," The inner $31\farcs5 \times 31\farcs5$ $600
+\times 600$ pc) region of the field is shown in more detail in Figure \ref{fig2}."
+ The emission [rom the nucleus is clearly extended. aloug the NW-SE direction., The emission from the nucleus is clearly extended along the NW–SE direction.
+ There are a total of 11 sources detected above a nominal So threshold within the central ~Ix1 re;[n])ion using tlie algoritliui (distributed as part of the data analysis package).," There are a total of 11 sources detected above a nominal $5\sigma$ threshold within the central $\simeq
+1^{\prime} \times 1^{\prime}$ region using the algorithm (distributed as part of the data analysis package)."
+ Table 1. gives the positious aud backgrouud subtracted count. rates for each source., Table \ref{tbl-1} gives the positions and background subtracted count rates for each source.
+ The background. was taken from soiwce [ree regious surrounding each source., The background was taken from source free regions surrounding each source.
+ The count rates are for the comined 0.ts aud 3.25 frame time data. except for sources D. EF. and J. which are allected by pileup in the longer frame time data.," The count rates are for the combined 0.4s and 3.2s frame time data, except for sources D, F, and J, which are affected by pileup in the longer frame time data."
+ We discuss the nuclear emission in Section 3.2.. the large scale extended emission in Section 3.3..) aud the compact uou-nuclear N-ray sources iu Section 3.L.," We discuss the nuclear emission in Section \ref{sec-nucleus}, the large scale extended emission in Section \ref{sec-large_scale}, and the compact non-nuclear X-ray sources in Section \ref{sec-compact}."
+ A radial profile of the X-ray surface briguuess was cetermiued by extracting. [rom the unpiled short frame time observation. couuts in circlar annuli with widths 1pixel (075) centered on the uucleus (Figure ) 3)).," A radial profile of the X-ray surface brightness was determined by extracting, from the unpiled short frame time observation, counts in circular annuli with widths $1~\rm pixel$ $0.\!^{\prime\prime}5$ ) centered on the nucleus (Figure \ref{fig3}) )."
+ The X-ray. emission [roiu compact sources close to the nucleus was excluded. aud the backerouucd was estimated [rom a €ircular region. 10” SW of the nucleus. which is free of large scale exteucded X-ray. emission.," The X-ray emission from compact sources close to the nucleus was excluded, and the background was estimated from a circular region $40^{\prime\prime}$ SW of the nucleus, which is free of large scale extended X-ray emission."
+ The telescope PSF at 2 keV was calculated using the program. aud file aciss1998-11-022dpst3N0002.its., The telescope PSF at 2 keV was calculated using the program and file aciss1998-11-052dpsf3N0002.fits.
+ We find the number of 0.5-8 keV counts within a radius of 3” is about twice that of the norm:ilizecl telescope PSF. within the same racliS.," We find the number of 0.5–8 keV counts within a radius of $3^{\prime\prime}$ is about twice that of the normalized telescope PSF, within the same radius."
+" The broad-baucl X-ray. euission [rom the nucleus shows a bright core plus au exteusion of c 3-41"" in position angle (p.a.)", The broad-band X-ray emission from the nucleus shows a bright core plus an extension of $\simeq 3$ $4^{\prime\prime}$ in position angle (p.a.)
+ ~290° (Figures l.. 2)).," $\simeq 290^{\circ}$ (Figures \ref{fig1}, \ref{fig2}) )."
+ This direction of X-ray extent is almost yerpendicular to that of the galaxy disk (p.a.=25°: Maiolino et al., This direction of X-ray extent is almost perpendicular to that of the galaxy disk $\rm p.a. = 25^{\circ}$; Maiolino et al.
+ 1998)., 1998).
+" optical images (Wilsonetal.2000) show a €olpact V-shaped structure (possibly au ""ionization cone"") extending ~2"" to the NW (in p.a.= 316°) of the nucleus."," optical emission-line images \citep{wil00} show a compact $\mathbf V$ -shaped structure (possibly an “ionization cone”) extending $\simeq 2^{\prime\prime}$ to the NW (in $\rm p.a. =
+316^{\circ}$ ) of the nucleus."
+ The direction of elongation of he nuclear X-ray. emission is close to that «X the bright southern edge of this 7V. which lies iu ya.9807.," The direction of elongation of the nuclear X-ray emission is close to that of the bright southern edge of this $\mathbf V$ ”, which lies in $\rm p.a. \simeq 280^{\circ}$."
+" Thus both the scale aud directici of the nuclear X-ray exteusion are close to those of he optical emission-line ""V.", Thus both the scale and direction of the nuclear X-ray extension are close to those of the optical emission-line $\mathbf V$ ”.
+ Observations of Circinus in water maser emission reveal gas both in a warped. edge-on accretion disk aud in outflow from the nucleus (Greenhill2001)..," Observations of Circinus in water maser emission reveal gas both in a warped, edge-on accretion disk and in outflow from the nucleus \citep{gre01a}."
+ The mean p.a., The mean p.a.
+ oL the maser outflow ts 3087. which is also quite close to thedirection oftheX-ray extension in the nucleus.," of the maser outflow is $308^{\circ}$ , which is also quite close to thedirection of theX-ray extension in the nucleus."
+ The distribution of Fe Ita line emitting gas was determined by extracting au nage in the 6.2-6.5 keVband. which iucludes all of the Fe Ina liue emission. aud subtracting the coutinuum," The distribution of Fe $\alpha$ line emitting gas was determined by extracting an image in the 6.2–6.5 keVband, which includes all of the Fe $\alpha$ line emission, and subtracting the continuum"
+"Black holes (DIIS) are very simple svsteius, cutirely characterized by mass and spin.","Black holes (BHs) are very simple systems, entirely characterized by mass and spin."
+ However. accretion onto DITs is a umch more complex process.," However, accretion onto BHs is a much more complex process."
+" Obscrvationally. accreting stellar-1nass DITs are found in various transientl states, with different spectral and timing properties, and a different balance between mass-enerev advection. radiation and outflows."," Observationally, accreting stellar-mass BHs are found in various transient states, with different spectral and timing properties, and a different balance between mass-energy advection, radiation and outflows."
+" For DIIs in a binary system, aceyotins from a Roche-lobe-filling donor star. the iain parameter that determines the accretion state js thought to be the mass accretion rate. or nore precisely the mass inflow rate M at the outer edee of the accretion disk."," For BHs in a binary system, accreting from a Roche-lobe-filling donor star, the main parameter that determines the accretion state is thought to be the mass accretion rate, or more precisely the mass inflow rate $\dot{M}$ at the outer edge of the accretion disk."
+" This can be larger than the accretion. rate through the BI horizon,. if. some of the is lost; in outflows."," This can be larger than the accretion rate through the BH horizon, if some of the gas is lost in outflows."
+; The inflow ↴∙rate depends on the and⋅ orbital separation⋅∙ of↴⋅∙ the companion star., The inflow rate depends on the nature and orbital separation of the companion star.
+dun nature It is convenient to rescale this rate as a cdimoensiouless quantity. 51= M/Mgq. where Afgaa=Lrgaa/(00.10) is. the accretion. rate that would be required. toproduce a huninositv L=πμ13«1075 cre + with a standard radiative efficiency = 0.1. .," It is convenient to rescale this rate as a dimensionless quantity, $\dot{m} \equiv \dot{M}/\dot{M}_{\rm Edd}$ , where $\dot{M}_{\rm Edd} = L_{\rm Edd}/(0.1c^2)$ is the accretion rate that would be required toproduce a luminosity $L = L_{\rm Edd} = 1.3 \times 10^{38}$ erg $^{-1}$ with a standard radiative efficiency $= 0.1$ ."
+ ↕≧↕↕↴∖↴⋜↕↸⊳↸⊳↥⋅↸∖↑∐↓∶↴∙⊾⋜↕↑⋜↕↥⋅⋜↕↑↸∖∣⊔⋞⊰⋜↕↕↸∖↖↖⇁↓∩−⋜∐⋅↸∖ ∙∕ . ↽⋅≻ eoncrally⋝ found. in: the low/hard; state (RemillardM aud AMeChutock.McClintock 2006).," BHs accreting at a rate $\dot{m} \la$ a few $10^{-2}$ are generally found in the low/hard state (Remillard and McClintock, 2006)."
+" AOObservationallv,rvatinally, this is cliaracterizeccharacterizedli bv a cool, faintστ thermal disk and a hard power-law component (photon iudex Pz1.5 2)."," Observationally, this is characterized by a cool, faint thermal disk and a hard power-law component (photon index $\Gamma \approx 1.5$ $2$ )."
+" The power-law enissonmiion «cones fromfro the ταradiativelvliatively inefficiinefficient inner region of the nieinflow, and is probably due to utiuverse-Compton-scattering iu a hot. opticallv-thiu corona or at the base of a jet."," The power-law emission comes from the radiatively inefficient inner region of the inflow, and is probably due to inverse-Compton-scattering in a hot, optically-thin corona or at the base of a jet."
+ The disk is sometimes truncated at a large distance from the imnermost stable circular orbit., The disk is sometimes truncated at a large distance from the innermost stable circular orbit.
+ Radio studies suggestOO (Fender et al..," Radio studies suggest (Fender et al.,"
+" 2001) that the power is carried mostout by a steady jet (detectedaccretion as Πατρούτα radio cliission), with a possible advected coniponent."," 2004) that most of the accretion power is carried out by a steady jet (detected as flat-spectrum radio emission), with a possible advected component."
+" When 0.1ZahX1. aecreting DIIS aretypically dominated by an opticallh-thick, ecometrically-thin"," When $0.1 \la \dot{m} \la 1$, accreting BHs aretypically dominated by an optically-thick, geometrically-thin"
+shown in Table 4 for all filters.,shown in Table \ref{tab:sunave} for all filters.
+" Besides reporting the zeropoint (mo, mag arcsec”), the slope (y, mag arcsec? ΜΙΥ|) and their associated statistical errors, the Table includes also the estimated full solar cycle variation (Am=|(2.4—0.8)y, the value attainedat solar minimum, evaluated for SFD=0.8 MJy (min), the value corresponding to the average SFD level <SFD>=1.6 MJy (m4), the RMS deviation from the best fit relation (c, the linear correlation factor (r) and the number of nights used (NV) for each filter."," Besides reporting the zeropoint $m_0$ , mag $^{-2}$ ), the slope $\gamma$, mag $^{-2}$ $^{-1}$ ) and their associated statistical errors, the Table includes also the estimated full solar cycle variation $\Delta
+m=|(2.4-0.8)\gamma|$ ), the value attainedat solar minimum, evaluated for $SFD$ =0.8 MJy $m_{min}$ ), the value corresponding to the average $SFD$ level $<SFD>$ =1.6 MJy $m_{ave}$ ), the RMS deviation from the best fit relation $\sigma$ ), the linear correlation factor $r$ ) and the number of nights used $N$ ) for each filter."
+" As one can see, the values of Am are smaller than those reported by other authors: with the only exception of U, which reaches about 0.6 mag arcsec?, all the others show values that are smaller than 0.4 mag arcsec?."," As one can see, the values of $\Delta m$ are smaller than those reported by other authors: with the only exception of $U$, which reaches about 0.6 mag $^{-2}$, all the others show values that are smaller than 0.4 mag $^{-2}$."
+" Walker (1988)) quoted maximum ranges of AV«1.0 and AB«0.8 mag arcsec? for solar cycle n. 21, while Krisciunas (1997)) reports AV=0.6 for solar cycle n. 22, and similar values are reported by Leinert et al. (1995))"," Walker \cite{walker88}) ) quoted maximum ranges of $\Delta V\simeq$ 1.0 and $\Delta B\simeq$ 0.8 mag $^{-2}$ for solar cycle n. 21, while Krisciunas \cite{krisc97}) ) reports $\Delta V$ =0.6 for solar cycle n. 22, and similar values are reported by Leinert et al. \cite{leinert95}) )"
+ and Mattila et al. (1996))., and Mattila et al. \cite{attila}) ).
+" On the other hand, Liu et al. (2003))"," On the other hand, Liu et al. \cite{liu}) )"
+ quote an increase of the V sky brightness of ~0.2 mag arcsec? from 1995 to 2001., quote an increase of the $V$ sky brightness of $\sim$ 0.2 mag $^{-2}$ from 1995 to 2001.
+" This value is consistent with the measures discussed here, especially taking into account that cycle n. 23 had a second maximum, which occurred after the observations presented by Liu et al. (2003))."," This value is consistent with the measures discussed here, especially taking into account that cycle n. 23 had a second maximum, which occurred after the observations presented by Liu et al. \cite{liu}) )."
+ These facts seem to suggest that not all solar cycles have identical effects on the night glow., These facts seem to suggest that not all solar cycles have identical effects on the night glow.
+" As a matter of fact, Walker (1988)), while revising the result of previous works, had suggested that the relation between intensity of the [OI]5577 line and the solar activity might vary from cycle to cycle, within a given cycle and possibly with geographical location."," As a matter of fact, Walker \cite{walker88}) ), while revising the result of previous works, had suggested that the relation between intensity of the [OI]5577 line and the solar activity might vary from cycle to cycle, within a given cycle and possibly with geographical location."
+" Unfortunately, the number of sunspot cycles covered by the observations is still too small to allow a firm conclusion, but the very recent results discussed by Krisciunas et al. (2007)),"," Unfortunately, the number of sunspot cycles covered by the observations is still too small to allow a firm conclusion, but the very recent results discussed by Krisciunas et al. \cite{krisc07}) ),"
+" covering two full solar cycles, seem indeed to confirm this suspicion."," covering two full solar cycles, seem indeed to confirm this suspicion."
+" Given the unprecedentedly large amount of data, one can investigate the solar dependency in a bit more detail, for example trying to deduce the typical timescales of night sky brightness fluctuations due to short-term changes in the solar flux."," Given the unprecedentedly large amount of data, one can investigate the solar dependency in a bit more detail, for example trying to deduce the typical timescales of night sky brightness fluctuations due to short-term changes in the solar flux."
+" Ideally, for doing this, one would look for strong enhancements in the solar flux and try to seek for a corresponding increases in the night sky brightness."," Ideally, for doing this, one would look for strong enhancements in the solar flux and try to seek for a corresponding increases in the night sky brightness."
+" Nevertheless, due to the sparse time sampling, the only viable approach is actually the opposite one, i.e. starting from the available night sky measurements, one goes back to the solar data and studies the correlation with the sun flux emitted as a function of time lag."," Nevertheless, due to the sparse time sampling, the only viable approach is actually the opposite one, i.e. starting from the available night sky measurements, one goes back to the solar data and studies the correlation with the sun flux emitted as a function of time lag."
+ For doing this I have computed the Pearson’s linear correlation coefficient r (Press et al. 1992)), For doing this I have computed the Pearson's linear correlation coefficient $r$ (Press et al. \cite{press}) )
+" between the nightly average sky brightness m(t) measured at any given time f and the solar flux density S—FD(t7), averaged within a time window Art, as a functionof both t and Art."," between the nightly average sky brightness $m(t)$ measured at any given time $t$ and the solar flux density $SFD(t-\tau)$, averaged within a time window $\Delta \tau$, as a functionof both $\tau$ and $\Delta
+\tau$ ."
+" While the first parameter gives an indication about the time lag between a change in the solar flux and the consequent variation in the night sky brightness, the second is related to the typical"," While the first parameter gives an indication about the time lag between a change in the solar flux and the consequent variation in the night sky brightness, the second is related to the typical"
+the observed sample.,the observed sample.
+ For a galaxy to be deected as a host at least one LGIUD must be detected. within it by a high-energy observatory., For a galaxy to be detected as a host at least one LGRB must be detected within it by a high-energy observatory.
+ Then. the galaxy itself must be detected. usually by optical/NII telescopes.," Then, the galaxy itself must be detected, usually by optical/NIR telescopes."
+ We woulc treat the biases introduced by each of these observations se»watelv., We would treat the biases introduced by each of these observations separately.
+ To model the first bias. we compute the probability of detecting at least one. LOGIUS within it by any of the high-energv observatories monitoring GRBs.," To model the first bias, we compute the probability of detecting at least one LGRB within it by any of the high-energy observatories monitoring GRBs."
+ Given that event detection. processes follow Poissonian statistics. if Ποια.) is the contribution of galaxy g at. redshilft to the LGRB rate observed by experiment 7. the probability of observing at least one LOGIUD in a galaxy is where AZ; is the time interval of the LOGRB search conducted. by experiment ὃν," Given that event detection processes follow Poissonian statistics, if $R_{{\rm GRB},i}(g,z)$ is the contribution of galaxy $g$ at redshift $z$ to the LGRB rate observed by experiment $i$, the probability of observing at least one LGRB in a galaxy is where $\Delta T_i$ is the time interval of the LGRB search conducted by experiment $i$."
+ Eqn., Eqn.
+ 5 assumes that the searches hy cilferent observatories are independent., \ref{phg} assumes that the searches by different observatories are independent.
+ Previous works estimate the intrinsic LGIUD rate in a tvpical galaxy to be of the order o£ 10.7vr.+ (Ervoretal.1999).. hence the LGRB rate observed in the galaxy by any mission must be lower than this value.," Previous works estimate the intrinsic LGRB rate in a typical galaxy to be of the order of $10^{-3}\,{\rm yr}^{-1}$ \citep{Fry99}, hence the LGRB rate observed in the galaxy by any mission must be lower than this value."
+ Given that current searches. [or LORBs have spannecl a few vears. »RonpjAL;m1. anc pualg.2)&»Ponga which means that the likelihood ofa galaxy being observed. as a host is proportional to its contribution to the meanobserced number of LOIDs. where AO; is the sky coverage of the experiment 7.," Given that current searches for LGRBs have spanned a few years, $\sum_i R_{{\rm GRB},i} \Delta T_i \ll 1$, and $p_{\rm HG}(g,z)
+\simeq \sum_i R_{{\rm GRB},i} \Delta T_i$, which means that the likelihood of a galaxy being observed as a host is proportional to its contribution to the mean number of LGRBs, where $\Delta \Omega_i$ is the sky coverage of the experiment $i$."
+ Note that. for a fixed. z. cosmological and observatory dependent [actors in Eqn.," Note that, for a fixed $z$, cosmological and observatory dependent factors in Eqn."
+ 6 become constant. independently. of the number of observatories considered. and cancel out in the normalization of the weights.," \ref{rateobsgal} become constant, independently of the number of observatories considered, and cancel out in the normalization of the weights."
+ This means that when computing mean LIC. properties such as mass or SER as a function of z. the first bias can be modeled simply. by weighting the corresponding properties of each galaxy by itsinbrinsiec LORB rate.," This means that when computing mean HG properties such as mass or SFR as a function of $z$, the first bias can be modeled simply by weighting the corresponding properties of each galaxy by its LGRB rate."
+ In the computation of the integrated. properties. of the whole population of host galaxies (like the integrated mass distribution. for example). also the effects of volume variation. time cdilation anc detectability of the cülferent observatories must be taken into account (Pellizza ct al.," In the computation of the integrated properties of the whole population of host galaxies (like the integrated mass distribution, for example), also the effects of volume variation, time dilation and detectability of the different observatories must be taken into account (Pellizza et al."
+ in preparation)., in preparation).
+ Ln this case it is better to compute first he mass clistribution observed by each experiment at cach redshift z. weighting the galaxies in the corresponding snapshot bv pgcot(g.z).," In this case it is better to compute first the mass distribution observed by each experiment at each redshift $z$, weighting the galaxies in the corresponding snapshot by $p_{\rm HG}(g,z)$."
+ Second. the integration in z can ο performed for each experiment. and the resulting distribution can be normalized.," Second, the integration in $z$ can be performed for each experiment, and the resulting distribution can be normalized."
+ This has the advantage of avoiding the use of the values of AY; and AQ;. which are »oorlv. known for some of the observatories that cletectoc he LORBs in the sample of Savagli," This has the advantage of avoiding the use of the values of $\Delta T_i$ and $\Delta \Omega_i$, which are poorly known for some of the observatories that detected the LGRBs in the sample of \cite{Sav09}."
+oetal.(2009).. Only ασ) for cach experiment is needed. which is computer as described in Sect.," Only $p_{{\rm det},i}(z)$ for each experiment is needed, which is computed as described in Sect."
+ 3.2. for pactwarsez).," \ref{obsrate} for $p_{\rm det,BATSE}(z)$."
+ Finally. these distributions can be combined into a single one by adding them. previously. scaled to the number of LORBs detectec bv cach observatory.," Finally, these distributions can be combined into a single one by adding them, previously scaled to the number of LGRBs detected by each observatory."
+ Ehe relevant cata for computing the cletectability were taken from Sternetal.(2001) for BATSE. Guetta&Piran(2007) forοι. Fronteraetal.(2009). forAX. Pélangeonetal.(2008). for ΗΕand Hurleyetal.(1992) forUlysses.," The relevant data for computing the detectability were taken from \citet{Ste01} for BATSE, \citet{Gue07} for, \citet{Fro09} for, \citet{Pel08} for and \citet{Hur92} for."
+ These experiments. 2detected: 36 of the 38 LORBs in the sample of Savaglioetal.(2009)., These experiments detected 36 of the 38 LGRBs in the sample of \citet{Sav09}.
+. bor statistical reasons. only experiments that detected. at least. five LGRBs in the sample were considered.," For statistical reasons, only experiments that detected at least five LGRBs in the sample were considered."
+Wind and NEAR-NGRS were discarded: because all their LORBs were observed also by other experiments above., and -XGRS were discarded because all their LGRBs were observed also by other experiments above.
+ To be consistent. the observed. mass distribution to which the model predictions were compared was constructed from the ata of Savaglioctal.(2009) bv adding the individual istributions observed by cach experiment.," To be consistent, the observed mass distribution to which the model predictions were compared was constructed from the data of \citet{Sav09} by adding the individual distributions observed by each experiment."
+ The second bias is more cillicult to. model., The second bias is more difficult to model.
+ iu the search [for host galaxies is guided. bv the iscovery of the LORBs themselves. anc usually done with a variety of cifferent telescopes ancl detectors. in ilferent. bands of the electromagnetic spectrum. ancl with ilerent sensitivities. the biases introduced: are unclear.," Given that the search for host galaxies is guided by the discovery of the LGRBs themselves and usually done with a variety of different telescopes and detectors, in different bands of the electromagnetic spectrum and with different sensitivities, the biases introduced are unclear."
+ ο.ralaxies with surface brightness below the detectability wesholel of available instruments could. introduce a bias owards low motallicity galaxies., Galaxies with surface brightness below the detectability threshold of available instruments could introduce a bias towards low metallicity galaxies.
+ An extra. problem cou of caused by cust) obscuration which could. allect the etection of the afterglows and produce a bias towards high =jetallicity galaxies (Evnboctal.2009)., An extra problem could be caused by dust obscuration which could affect the detection of the afterglows and produce a bias towards high metallicity galaxies \citep{Fyn09}.
+. Hence. at. leas rese two elfects might combine themselves to. determine 1e cletectability of a host galaxy.," Hence, at least these two effects might combine themselves to determine the detectability of a host galaxy."
+ Considering also the fac jiu we have only access to the public galaxy catalogue of ο which provides only mean global properties ancl magnitudes. we adopt the observec integrated stellar mass distribution as a tool to apply the combined ellect of observational biases to the. simulate galaxy sample.," Considering also the fact that we have only access to the public galaxy catalogue of the which provides only mean global properties and magnitudes, we adopt the observed integrated stellar mass distribution as a tool to apply the combined effect of observational biases to the simulated galaxy sample."
+ We use the fact that the observed integratec stellar mass distribution has been allected by observationa Xases although. we cannot disentanegle their. individua effects., We use the fact that the observed integrated stellar mass distribution has been affected by observational biases although we cannot disentangle their individual effects.
+" οσο, we require the simulated. LGIUD. hosts to reproduce the observed stellar mass. distribution in order o build up the observable simulated: LGORB hosts."," Hence, we require the simulated LGRB hosts to reproduce the observed stellar mass distribution in order to build up the observable simulated LGRB hosts."
+ The »ocedure is explained in detail in next section., The procedure is explained in detail in next section.
+ Note that. jereafterin. we will discuss the trends of the observable simulated. hosts. and the general. galaxy population.," Note that, hereafterin, we will discuss the trends of the observable simulated hosts and the general galaxy population."
+ The ormer can be compared to the current observed hosts. out. if this later sample changes due to better or cillerent observational techniques. our simulated sample should. be also consistently modified to match the new observed stellar mass cistribution.," The former can be compared to the current observed hosts, but, if this later sample changes due to better or different observational techniques, our simulated sample should be also consistently modified to match the new observed stellar mass distribution."
+ The largest. ancl most. comprehensively uniform sample of rost galaxy properties available at. present is that compiled ov Savaglioctal.(2009)erbhosts., The largest and most comprehensively uniform sample of host galaxy properties available at present is that compiled by \citet{Sav09}.
+info.. Stellar masses. star formation rates. metallicities. absolute magnitudes and colours. of 46 observed. host. galaxies. up O0 23 were obtained by these authors comparing their spectral energy distributions to those of svnthetie stellar »opulations. using the method described in (2004).," Stellar masses, star formation rates, metallicities, absolute magnitudes and colours of 46 observed host galaxies up to $z \sim 3$ were obtained by these authors comparing their spectral energy distributions to those of synthetic stellar populations, using the method described in \citet{Gla04}."
+. A Κον point of our mocels is that we require the stellar mass distribution of the predicted. host. galaxies to match hat of the observed host galaxies., A key point of our models is that we require the stellar mass distribution of the predicted host galaxies to match that of the observed host galaxies.
+ The stellar. mass is adopted as the property to be reproduced by. the models since it is now widely accepted that stellar. mass is a more fundamental quantity for galaxies than luminosity, The stellar mass is adopted as the property to be reproduced by the models since it is now widely accepted that stellar mass is a more fundamental quantity for galaxies than luminosity
+is shown in figure 3.,is shown in figure 3.
+ At the high mass eud he sources nav be very Wmiuinous aud cau thus ve detected to laree distances. whereas only the closest of the low mass mini-halos are detectable.," At the high mass end the sources may be very luminous and can thus be detected to large distances, whereas only the closest of the low mass mini-halos are detectable."
+ The competition between increasing iminihlalo iuuber density aud decliniug Iuninosityv. towards ower miui-halo masses. results iu a characteristic (anediau) mass of 5\109NL. for the detectable sources.," The competition between increasing mini-halo number density and declining luminosity, towards lower mini-halo masses, results in a characteristic (median) mass of $5\times10^6\,{\rm M_\odot}$ for the detectable sources."
+ In the vicinity of the Sum. a iuni-halo of this mass is expected to have a luuinesity of approximately 8<Loleres1. du photons of energv above 100 MeV. This result i9 in accord with the analysis of the UID ECRET source population by Ianbach et al (1996). who couclided that them typical huunmositv is ~LO?eres+.," In the vicinity of the Sun, a mini-halo of this mass is expected to have a luminosity of approximately $8\times10^{34}\,{\rm erg\,s^{-1}}$, in photons of energy above 100 MeV. This result is in accord with the analysis of the UID EGRET source population by Kanbach et al (1996), who concluded that their typical luminosity is $\sim10^{35}\,{\rm erg\,s^{-1}}$."
+ The median line-of-sielit distance for the müuihalos with flux greater than the EGRET detection Init is 3.9 kpc., The median line-of-sight distance for the mini-halos with flux greater than the EGRET detection limit is 3.9 kpc.
+ Bearing in ind that both Active Galactic Nuclei (AGN) and pulsars ave effectively. poiut-like eammacray sources. the nowtrivial angular size predicted for the gamma-ray cluission frou dark uiini-halos is an duportaut feature of the model.," Bearing in mind that both Active Galactic Nuclei (AGN) and pulsars are effectively point-like gamma-ray sources, the non-trivial angular size predicted for the gamma-ray emission from dark mini-halos is an important feature of the model."
+" We characterise the size of the sources by 0,. the tidal radius of the mini-halo: this is he radius which cucloses all of the flix from the iiodcl source."," We characterise the size of the sources by $\theta_t$, the tidal radius of the mini-halo; this is the radius which encloses all of the flux from the model source."
+" We compute 0, using the same procedure as WOALO3.", We compute $\theta_t$ using the same procedure as WOM03.
+ For a singular isothermal sphere. the surface brightness is fx1/0 at all radii. aud we expect roughlv half of the flux to be contained within 0/2.," For a singular isothermal sphere, the surface brightness is $I\propto1/\theta$ at all radii, and we expect roughly half of the flux to be contained within $\theta_t/2$ ."
+" The median value of the tidal radius js (),;= 2.2"". for sources detectable by EGRET."," The median value of the tidal radius is $\langle\theta_t\rangle=2.2^\circ$ , for sources detectable by EGRET."
+ The full distribution of 0; is shown in figure L., The full distribution of $\theta_t$ is shown in figure 4.
+ Tutrinsic source confusion. ie. the regine where sources are so nunerous that they overlap on the sky. becomes a nünor problem iu the ↥⋅↸∖∶↴∙⊾↕∪↕⋜∐⋅∪∏∐≼⇂∐↸∖≼∶⋜↧↕⋯⊳↑↕↸⊳≼," Intrinsic source confusion, i.e. the regime where sources are so numerous that they overlap on the sky, becomes a minor problem in the region around the Galactic Centre."
+⊲↸∖∐⊓⋅↸∖∙∖∏↑∐↕∐∐∣↴ latitude/longitude of the Galactic Ceutre (an area of roughly 100 deg?). we find 32 sources; covering about 100deg? (the sources are smaller than average).," Within $10^\circ$ latitude/longitude of the Galactic Centre (an area of roughly $400\;{\rm deg^2}$ ), we find 32 sources, covering about $100\;{\rm deg^2}$ (the sources are smaller than average)."
+ This fure is an overestimate im the seuse that our model predicts too many sources see 866., This figure is an overestimate in the sense that our model predicts too many sources — see 6.
+" With more sensitive instiriuneutation. having a lower flux limit. intriusic source coufusion becomes a luajor issn, 1iini-halo population covers the entire sky (VONMastheG3)."," With more sensitive instrumentation, having a lower flux limit, intrinsic source confusion becomes a major issue, as the mini-halo population covers the entire sky (WOM03)."
+ It is clear that the model we have preseuted predicts too mauv ECRET sources., It is clear that the model we have presented predicts too many EGRET sources.
+ The cutire EGRET source list contaius only about as many sources as our simulation predicts. and παν (C405€)) of the cataloened sources are considered o be identified.," The entire EGRET source list contains only about as many sources as our simulation predicts, and many $\sim40$ ) of the catalogued sources are considered to be identified."
+" Furthermore. we nmst expect at there will be many examples of known types of σαΑΝ sources, such as pulsus. amongst rw UID. EGRET sources. plying that our uodel over-prediets the source count bv a factor zo2."," Furthermore, we must expect that there will be many examples of known types of gamma-ray sources, such as pulsars, amongst the UID EGRET sources, implying that our model over-predicts the source count by a factor $\ga2$."
+ However. it must be acknowledged that 1e inodoel we have presented involves substantial extrapolation iuto poorly charted territory aud is verefore illustrative. uot definitive.," However, it must be acknowledged that the model we have presented involves substantial extrapolation into poorly charted territory and is therefore illustrative, not definitive."
+ There are. or exunple. considerable uncertainties in the cosndc-crav distribution model. in the shape of je Galaxy's dark halo. aud in the mass spectrum of the clustering.," There are, for example, considerable uncertainties in the cosmic-ray distribution model, in the shape of the Galaxy's dark halo, and in the mass spectrum of the clustering."
+" Iu addition. we can mauediatelv )mt to one aspect of the dark matter model which leads to an over-prediction of the source counts: we have assumed that of the dark. uatter in the Galactic halo is bound iuto iiu (822: WOAL03), but this cannot be the case."," In addition, we can immediately point to one aspect of the dark matter model which leads to an over-prediction of the source counts: we have assumed that of the dark matter in the Galactic halo is bound into mini-halos 2; WOM03), but this cannot be the case."
+ Even if all the dark matter were nütilly in nuui-talos. the process of tidal stripping will gradually uubind material. leading to a significant fraction in he form of dark matter “streams”.," Even if all the dark matter were initially in mini-halos, the process of tidal stripping will gradually unbind material, leading to a significant fraction in the form of dark matter “streams”."
+ Considering hese uncertiüuties we recognise that the overποιο of source counts is uot a fundamental deficicney of the model.," Considering these uncertainties, we recognise that the over-prediction of source counts is not a fundamental deficiency of the model."
+ Other basic features of the UID source population can be used to gauge the success of the model: the distribution on the sky: spectra: variability: angular structure: aud observations made in other wave-hauds., Other basic features of the UID source population can be used to gauge the success of the model: the distribution on the sky; spectra; variability; angular structure; and observations made in other wave-bands.
+ We now address cach of these aspects in turn., We now address each of these aspects in turn.
+ The model correctly predicts a prepouderance of sources at low- to mid-Galactic latitudes. with about the right median latitude (127) for the sopulation.," The model correctly predicts a preponderance of sources at low- to mid-Galactic latitudes, with about the right median latitude $12^\circ$ ) for the population."
+" Tlowever., it does uot exhibit the wo component (thindisk plas thick-disk) source »opulation which the data suggest (Golirels et al 2001),"," However, it does not exhibit the two component (thin-disk plus thick-disk) source population which the data suggest (Gehrels et al 2001)."
+ Fienre 2 shows a relatively strong vonceutration ∢ ⋅∢ ∐↸∖≼↧⋜↧⋜↧∙∐∪↖↖↽↸∖↖↽↸∖↥⋅∙↕∐↑∐↴∖↴↥⋅↸∖∶↴∙⊾↕," Figure 2 shows a relatively strong concentration towards the Galactic Centre region, compared to the data."
+∪∐↑↕∐∖↻↥⋅↸∖≼∐↸⊳↑↸∖≼↧⋅⋅↜⋅ ⋅ ⋅ source deusity is sufficiently lueh that source confusion isexpectedto be a major problem with observatious of∙ a∙ such a population.," However, in this region the predicted source density is sufficiently high that source confusion isexpectedto be a major problem with EGRET observations of a such a population."
+ The, The
+ TheE, The
+ TheEC, The
+ TheECR, The
+ TheECRE, The
+ TheECRET, The
+indicated in Figure 6..,indicated in Figure \ref{fig:n51d_image}.
+" Although a small circular region around the Wolt-Ravet star TD 36102 was excluded frou, the spectral analysis in the previous work. here we meluded this region because uo local structures were resolved in the low-resolution NIS. nuage."," Although a small circular region around the Wolf-Rayet star HD 36402 was excluded from the spectral analysis in the previous work, here we included this region because no local structures were resolved in the low-resolution XIS image."
+ The BCD reeion was selected to be an off-source polveoual region. but the CCD corner irradiated by the calibration sources was excluded.," The BGD region was selected to be an off-source polygonal region, but the CCD corner irradiated by the calibration sources was excluded."
+ Using4. the EPIC image (Figure 6ec). we searched for hard discrete sources.," Using the EPIC image (Figure \ref{fig:n51d_image}c c), we searched for hard discrete sources."
+ As a result. ten sources were detected in the SRC region. while uo point source was found in the BOD region.," As a result, ten sources were detected in the SRC region, while no point source was found in the BGD region."
+ These sources were not clearly resolved in the NIS nuage., These sources were not clearly resolved in the XIS image.
+ Therefore. we first analyzed the EPIC spectra of the point sources aud then performed a imodel fit to the NIS data in the SRC reeiou by adding the poiut-source commpoucut as a fixed model.," Therefore, we first analyzed the EPIC spectra of the point sources and then performed a model fit to the XIS data in the SRC region by adding the point-source component as a fixed model."
+ EPIC spectra of the teu poiut sources iu the SRC region were extracted aud integrated., EPIC spectra of the ten point sources in the SRC region were extracted and integrated.
+ The backgrouud spectra were extracted from an annular region around cach point source., The background spectra were extracted from an annular region around each point source.
+ The resultant backeround-subtracted spectra are shown in Figure 7.., The resultant background-subtracted spectra are shown in Figure \ref{fig:n51d_ptsrc}.
+ The spectra were well ⋅ ⋅ ↥⋅↸∖↻↥⋅∪≼↧⋯⊳↸∖≼↧↖∏↑∐⋜⋯⋜∏⋝↴∖↴∪∏≻⊓∪∐↻∪↖↖⇁↸∖↥⋅≓↕⋜∏↖↽⋯∪≼∐∖↕≺∖−≺↧∪↕ == Bl1/91)., The spectra were well reproduced with an absorption power-law model = 81/94).
+ The best-&t absorption column aud photon iudex were == 1.3 (0.87LS) 1075 om? and D = L3 (1.6LA). respectivelv.," The best-fit absorption column and photon index were = 1.3 (0.87–1.8) $\times 10^{21}$ $^{-2}$ and $\Gamma$ = 1.7 (1.6–1.8), respectively."
+ The observed flux was obtained to be 2.1 (2.32.5) «10P? in the 210 keV baud., The observed flux was obtained to be 2.4 (2.3–2.5) $\times 10^{-13}$ in the 2–10 keV band.
+ The fiux distribution of the teu sources is given in Table , The flux distribution of the ten sources is given in Table \ref{tab:ptsrc_list}.
+The spectral analysis of theSuzaku data was performed by essentially the same procedure as that described iu Section 3.3.., The spectral analysis of the data was performed by essentially the same procedure as that described in Section \ref{ssec:n11_suzaku}.
+ An example of a raw spectrum in tle SRC reeion and the corresponding NXD are compared in Figure d((a)., An example of a raw spectrum in the SRC region and the corresponding NXB are compared in Figure \ref{fig:n51d_compare}( (a).
+ Caven that both spectra exhibit the same intensity of cach fluorescent liue. we considered that the NNB was successfully constructed.," Given that both spectra exhibit the same intensity of each fluorescent line, we considered that the NXB was successfully constructed."
+ The NNB-subtracted spectra aud the 2.10 keV surface brightuesses of the SRC and BOD regions were obtained as shown in Figure 8((b) and Table 5.. respectively.," The NXB-subtracted spectra and the 2–10 keV surface brightnesses of the SRC and BGD regions were obtained as shown in Figure \ref{fig:n51d_compare}( (b) and Table \ref{tab:sbn_n51d}, respectively."
+ The SRC spectrum appears to show excess surface brightness iu the hard N-rav laud., The SRC spectrum appears to show excess surface brightness in the hard X-ray band.
+ However. upon subtracting the poiut-source flux from the SRC. we found that the remaining surface brightuess is consistent with that in the DGD region.," However, upon subtracting the point-source flux from the SRC, we found that the remaining surface brightness is consistent with that in the BGD region."
+ We applied the vieuctting correction to the BCD spectra and then subtracted them from tle SRC spectra., We applied the vignetting correction to the BGD spectra and then subtracted them from the SRC spectra.
+ The correction factors ΓΕ} were obtained as given im Figure 9.., The correction factors $f(E)$ were obtained as given in Figure \ref{fig:n51d_ratio}.
+ The resultant SRC spectra are shown in Figure 10.., The resultant SRC spectra are shown in Figure \ref{fig:n51d_diffuse}.
+ The spectra of two active FIs were merged to improve the photon statistics., The spectra of two active FIs were merged to improve the photon statistics.
+ Several cussion lines can be observed below 1 keV. while the spectra iu the hard band is featureless.," Several emission lines can be observed below $\sim$ 1 keV, while the spectrum in the hard band is featureless."
+ Therefore. we fitted the," Therefore, we fitted the"
+rez AMelénnedez (2005b). with the possible exception of DV which remains somewhat controversial.,"rez Melénndez (2005b), with the possible exception of $B-V$ which remains somewhat controversial."
+ While in the past the discrepancy was imputed to the temperature scale (Sekiguchi Fukuegita 2000). this is now no longer the issue.," While in the past the discrepancy was imputed to the temperature scale (Sekiguchi Fukugita 2000), this is now no longer the issue."
+ The main dillerence seems to be between studies deriving (D.V).~0.62 from fitting Yip colour relations from a wide range of temperatures and metallicities (Sekiguchi Fukugita 2000: rez Alelénndez 2005b). ancl on the other side studies focussing on solar analogues preselected. spectroscopicallv or within a narrower range of solarlike temperatures (Cavrel de Strobel 1996: Cray 1995: Soubiran. Triaucl 2004: present work) which favour (23).~0.64.," The main difference seems to be between studies deriving $(B-V)_{\odot} \sim 0.62$ from fitting $T_{\rm eff}$ –colour relations from a wide range of temperatures and metallicities (Sekiguchi Fukugita 2000; rez Melénndez 2005b), and on the other side studies focussing on solar analogues pre–selected spectroscopically or within a narrower range of solar–like temperatures (Cayrel de Strobel 1996; Gray 1995; Soubiran Triaud 2004; present work) which favour $(B-V)_{\odot} \sim 0.64$."
+ Increased accuracy in selecting good quality. homogeneous data has also some ellect on our result.," Increased accuracy in selecting good quality, homogeneous data has also some effect on our result."
+ Our (D.V).=0.642 is in close agreement with the colours of the closest. solar twin known so far (LR GOGO) and of the new candidate 11 178399., Our $(B-V)_{\odot}=0.642$ is in close agreement with the colours of the closest solar twin known so far (HR 6060) and of the new candidate HIP 78399.
+ Finally. once errors are properly taken into account. we remark that there is no crucial discrepaney between our result and the bluer colour of Sekiguchi Pulkugita. rez Melénndez.," Finally, once errors are properly taken into account, we remark that there is no crucial discrepancy between our result and the bluer colour of Sekiguchi Fukugita, rez Melénndez."
+ The main limit presently is the small number of stars with good quality metallicity data. having suitable CBVRL photometry.," The main limit presently is the small number of stars with good quality metallicity data, having suitable $UBVRI$ photometry."
+ We plan to improve on this aspect with new observations in the near future., We plan to improve on this aspect with new observations in the near future.
+ We thank Luca Casagrande for valuable advice on the HUPM and on the theoretical solar colours: and Ivan rez for useful correspondence., We thank Luca Casagrande for valuable advice on the IRFM and on the theoretical solar colours; and Ivan rez for useful correspondence.
+ This research. was supported by the Academy of Finland (grants nr., This research was supported by the Academy of Finland (grants nr.
+ 206055 and 208792)., 206055 and 208792).
+ This publication makes use of data products from the Two Micron. All Sky Survey. which is à joint project. of. the University of Massachusetts and the Enfrared Processing and Analysis Center/Calilornia Institute of Technology. funded by the National Aeronautics and Space AXcdministration and the National Science Foundation.," This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
+but still considerably lower than the adfs of ~700 seen in the polar knots of Abell 30 (Wesson et al.,but still considerably lower than the adfs of $\sim$ 700 seen in the polar knots of Abell 30 (Wesson et al.
+ 2003)., 2003).
+ Abundances for 5358s knot were previously derived by GM96., Abundances for 58's knot were previously derived by GM96.
+ They adopted a temperature of 4450KK and c(11:7)20.29. and from CELs they derived value of O about eight times lower than our value and a value of — about three times higher.," They adopted a temperature of K and $c({\rm H}\beta)$ =0.29, and from CELs they derived value of $^{2+}$ $^+$ about eight times lower than our value and a value of $^+$ $^+$ about three times higher."
+ They suggested that oxygen abundances were unreliable due to shock excitation of the knot., They suggested that oxygen abundances were unreliable due to shock excitation of the knot.
+ However. we calculate the ratio of energy input from shocks and from photoionisation as follows: the luminosity of the central star is [O!L.;. and its effective temperature is 0000KK (Clayton et al 2006).," However, we calculate the ratio of energy input from shocks and from photoionisation as follows: the luminosity of the central star is $^4$ $_{\sun}$, and its effective temperature is K (Clayton et al 2006)."
+ Based on model atmospheres of hydrogen-deticient stars by Rauch et al. , Based on model atmospheres of hydrogen-deficient stars by Rauch et al. (
+003). we estimate that of the energy from the star is emitted shortward of912A. giving a total ionising luminosity of 10°L.:.,"2003), we estimate that of the energy from the star is emitted shortward of, giving a total ionising luminosity of $\times$ $^3$ $_{\sun}$."
+ The impact of the stellar wind on the knot will be much more significant than the expansion of the knot into the low-density outer nebula in producing shocks., The impact of the stellar wind on the knot will be much more significant than the expansion of the knot into the low-density outer nebula in producing shocks.
+ Adopting M -2-13.10 ΜΜ. toand v4 = + from Clayton et al (2006). we calculate that the wind luminosity is only 67L.;. or of the ionising luminosity.," Adopting Ṁ = $\times$ $^{-7}$ $_{\sun}$ $^{-1}$ and $_{\infty}$ = $^{-1}$ from Clayton et al (2006), we calculate that the wind luminosity is only $_{\sun}$, or of the ionising luminosity."
+ Therefore photoionisiation dominates over shock excitation in producing the observed spectrum., Therefore photoionisiation dominates over shock excitation in producing the observed spectrum.
+ GM96 found OF = 0.05. while we obtain a value of ~ 1.," GM96 found $^{2+}$ $^{+}$ = 0.05, while we obtain a value of $\sim$ 1."
+" This difference is partly accounted for by the lower extinction and temperature adopted by GM96: we derived abundances from their observed fluxes but using c(117)22.0 and. 7,220 KK. and obtained O. /O = —0.4."," This difference is partly accounted for by the lower extinction and temperature adopted by GM96: we derived abundances from their observed fluxes but using $c({\rm H}\beta)$ =2.0 and $T_{\rm e}$ K, and obtained $^{2+}$ $^{+}$ = $\sim$ 0.4."
+ The remaining factor of 2.5 may be due to the temperature of the central star. and thus the ionisation degree of the nebula. increasing during the nine years between GM96's observations and our own.," The remaining factor of 2.5 may be due to the temperature of the central star, and thus the ionisation degree of the nebula, increasing during the nine years between GM96's observations and our own."
+ Table 8 presents a comparison between the mass fractions of H. He. C. N. O and Ne that we have derived for the ejecta of A 58 and A 30. and those derived by other authors for the central stars of these two objects. as well as for a number of other objects. including PG 1159-035. which may be a descendant of a star similar to the nucleus of A 30: V4334 Ser which has been proposed to be a younger counterpart to A 58/V605 Aqgl: the shells of two old novae. DQ Her and RR Pic: and the ejecta of three neon novae.," Table \ref{comparisons} presents a comparison between the mass fractions of H, He, C, N, O and Ne that we have derived for the ejecta of A 58 and A 30, and those derived by other authors for the central stars of these two objects, as well as for a number of other objects, including PG 1159-035, which may be a descendant of a star similar to the nucleus of A 30; V4334 Sgr, which has been proposed to be a younger counterpart to A 58/V605 Aql; the shells of two old novae, DQ Her and RR Pic; and the ejecta of three neon novae."
+ The nitrogen and neon abundances listed for the A 58 central knot are given in parentheses. since they are based on its measured CEL N:O:Ne ratios. normalised to its ORL oxygen abundance.," The nitrogen and neon abundances listed for the A 58 central knot are given in parentheses, since they are based on its measured CEL N:O:Ne ratios, normalised to its ORL oxygen abundance."
+ This analysis of the optical spectrum of 558 reveals many similarities with the much older knots of 330., This analysis of the optical spectrum of 58 reveals many similarities with the much older knots of 30.
+ In both cases. extreme abundance discrepancy factors are observed. while ORL temperature diagnostics indicate the existence of very cold ionised plasma.," In both cases, extreme abundance discrepancy factors are observed, while ORL temperature diagnostics indicate the existence of very cold ionised plasma."
+ Modelling of 330 by Ercolano et al. (, Modelling of 30 by Ercolano et al. (
+2003) demonstrated the physical plausibility of this situation. with the very high abundances of CNO coolants resulting in à very cold ionised core. with dust photoelectric heating generating the very high [O ΠΠ temperatures in the outer part of the knot.,"2003) demonstrated the physical plausibility of this situation, with the very high abundances of CNO coolants resulting in a very cold ionised core, with dust photoelectric heating generating the very high [O ] temperatures in the outer part of the knot."
+ The C/O number ratio in both knots of A330 was found to be less than unity (~0.3 for both knots when derived from ORLs. and 0.8 for knot J3 from CELs). contrary to the predictions of the “born-again scenario (Iben et al.," The C/O number ratio in both knots of 30 was found to be less than unity $\sim$ 0.3 for both knots when derived from ORLs, and 0.8 for knot J3 from CELs), contrary to the predictions of the `born-again' scenario (Iben et al."
+ 1983. Herwig 2001).," 1983, Herwig 2001)."
+ The knot of 558 is also found to be oxygen-rich. with a C/O ratio derived from ORLs of just 0.09.," The knot of 58 is also found to be oxygen-rich, with a C/O ratio derived from ORLs of just 0.09."
+ The single-star VETP model has serious difficulties in this scenario., The single-star VLTP model has serious difficulties in this scenario.
+ It is noticeable that the C/O mass fraction ratios measured for the A 58 and A 30 knots. 1/11 and 1448 respectively. are much smaller than the ratios of of 8 and 2.7 measured for their central stars.," It is noticeable that the C/O mass fraction ratios measured for the A 58 and A 30 knots, 1/11 and 1/4 respectively, are much smaller than the ratios of of 8 and 2.7 measured for their central stars."
+ In addition. the A 58 and A 30 knots have very large neon mass fractions.," In addition, the A 58 and A 30 knots have very large neon mass fractions."
+ We conclude that the signiticant differences between the C. O and Ne mass fractions measured for the knots and for the central stars are real.," We conclude that the significant differences between the C, O and Ne mass fractions measured for the knots and for the central stars are real."
+ Wesson et al. (, Wesson et al. (
+2003) pointed out that the C/O «1 ratios that they found for the knots of A 30 were not predicted by VLTP models for born-again nebulae.,2003) pointed out that the C/O $< 1$ ratios that they found for the knots of A 30 were not predicted by VLTP models for born-again nebulae.
+ Werner Herwig (2006) have recently reviewed VLTP models. confirming that C/O ratios greater than unity are predicted. consistent with the central star C/O ratios but not the ejecta C/O ratios found for A 58 and A 30.," Werner Herwig (2006) have recently reviewed VLTP models, confirming that C/O ratios greater than unity are predicted, consistent with the central star C/O ratios but not the ejecta C/O ratios found for A 58 and A 30."
+ In addition. VLTP models predict neon mass fractions 3:0.02. consistent with the values found for WCE central stars and PG 1159 stars. but much smaller than the neon mass fractions found in the knots of A 58 and A 30.," In addition, VLTP models predict neon mass fractions $\leq$ 0.02, consistent with the values found for WCE central stars and PG 1159 stars, but much smaller than the neon mass fractions found in the knots of A 58 and A 30."
+ The 558 knot is presumed to have a red-shifted counterpart. on the opposite side of the star and completely obseured by dust.," The 58 knot is presumed to have a red-shifted counterpart, on the opposite side of the star and completely obscured by dust."
+ The collimation of 330s knots argues against a single star being their source. and if 558 also has collimated outflows. as suggested by Pollacco et al. (," The collimation of 30's knots argues against a single star being their source, and if 58 also has collimated outflows, as suggested by Pollacco et al. ("
+1992). this could argue for some mechanism involving an aecretion dise and polar jets in a double star system.,"1992), this could argue for some mechanism involving an accretion disc and polar jets in a double star system."
+ In contrast to 330. though. the knot of 558 is expanding quite," In contrast to 30, though, the knot of 58 is expanding quite"
+"aare from wide-field IRAS surveys (Lonsdale et al.,","are from wide-field IRAS surveys (Lonsdale et al.,"
+" 1990; Rowan-Robinson et al.,"," 1990; Rowan-Robinson et al.,"
+" 1991; Gregorich et al.,"," 1991; Gregorich et al.,"
+" 1995), aand ddata come from more recent deep surveys (Dole et al,"," 1995), and data come from more recent deep surveys (Dole et al.,"
+" 2004; Bétthermin et aL,"," 2004; Bétthermin et al.,"
+" 2010) and counts at 500umcome from BLAST (Patanchon et al.,"," 2010) and counts at come from BLAST (Patanchon et al.,"
+" 2009) and (Oliver et al.,"," 2009) and (Oliver et al.,"
+" 2010; Dunlop et al.,"," 2010; Dunlop et al.,"
+ 2010) observations., 2010) observations.
+" In the far-IR (upper three panels) there appears to be a continuation of the result we found at24um:: the galaxy model (solid black line), which continues to be dominated by the Mrk33 starburst component (dot-dashed black line), matches the bright end of the counts well, but perhaps begins to underpredict the faint end — this can be seen in the faintest ccounts."," In the far-IR (upper three panels) there appears to be a continuation of the result we found at: the galaxy model (solid black line), which continues to be dominated by the Mrk33 starburst component (dot-dashed black line), matches the bright end of the counts well, but perhaps begins to underpredict the faint end – this can be seen in the faintest counts."
+" Our result at lis consistent with that of Busswell Shanks (2001), who also achieved a good fit to the ccounts using a similar model."," Our result at is consistent with that of Busswell Shanks (2001), who also achieved a good fit to the counts using a similar model."
+ While the fit to the, While the fit to the
+"where 0.—ΚΙ fmc. nds the electron (ion)number density. σ the Thomson cross section. or the fine structure constant. and i, the electron mass.","where $\theta_e \equiv k T_e / m_e c^2$ , $n$ is the electron (ion)number density, $\sigma_T$ the Thomson cross section, $\alpha_f$ the fine structure constant, and $m_e$ the electron mass."
+ In this work. we assume pure hvdrogen gas.," In this work, we assume pure hydrogen gas."
+ The emitted bremsstralilung photons are upscattered by inverse Compton scattering., The emitted bremsstrahlung photons are upscattered by inverse Compton scattering.
+" The amplification of photon energy by single scattering is 1=12-40,+1667.", The amplification of photon energy by single scattering is $A \equiv 1+4\theta_e+16\theta_e^2$.
+ The probability of single scattering is P?=1—exp(—7.) (see e.g.. Dermer et al.," The probability of single scattering is $P = 1 - \exp(-\tau_{es})$ (see e.g., Dermer et al."
+ 1991)., 1991).
+ Therefore the mean amplification factor by Compton scattering would be The fraction 1—P of photons remain unscattered while (he fraction /? of photons are upscattered to ;1 times their initial energy., Therefore the mean amplification factor by Compton scattering would be The fraction $1-P$ of photons remain unscattered while the fraction $P$ of photons are upscattered to $A$ times their initial energy.
+ This prescription is valid only [or a single scaltering Comptonization., This prescription is valid only for a single scattering Comptonization.
+ Since most CDAF [flow considered here has τε<1. we use this amplification [actor lor Comptonized bremsstrahling.," Since most CDAF flow considered here has $\tau_{es} < 1$, we use this amplification factor for Comptonized bremsstrahlung."
+ Although Dermer et al. (, Although Dermer et al. (
+"1991) provides a handy formula For η, applicable to diverse regimes. il can give incorrect values for Am 1and naQ<1. the main parameter regime of CDAF.","1991) provides a handy formula for $\eta$ , applicable to diverse regimes, it can give incorrect values for $A \gg 1$ and $\tau_{es}
+\lesssim 1$, the main parameter regime of CDAF."
+" Although τι is a hard-to-deltine quantity in a complex flow. we simply adopt 7,=n,oyr in Chis work for the purpose of estimating the Comptonization."," Although $\tau_{es}$ is a hard-to-define quantity in a complex flow, we simply adopt $\tau_{es}
+= n_e \sigma_T r$ in this work for the purpose of estimating the Comptonization."
+" When Pelx 1. Comptonization becomes saturated. and all photons that are not absorbed are upscattered to 37). that is 7—tar=36,/.7 where ic=Γι (Dermer et al."," When $PA \gg 1$ , Comptonization becomes saturated, and all photons that are not absorbed are upscattered to $3T_e$, that is $\eta \rightarrow
+\eta_{sat} = 3\theta_e /x$ where $x \equiv h\nu /m_e c^2$ (Dermer et al."
+ 1991)., 1991).
+ Then theFully saturated Comptonized bremsstrahlung emissionis, Then thefully saturated Comptonized bremsstrahlung emissionis
+v= [generatious.,$n=4$ generations.
+ The hatched bar iu panels e aud 5 shows he fraction of primordial gas remaining. as deteruuned x Q.," The hatched bar in panels $a$ and $b$ shows the fraction of primordial gas remaining, as determined by $Q$."
+ The analytic curves show the correspouding results roni equations 6 and 8.., The analytic curves show the corresponding results from equations \ref{nz} and \ref{ntot}.
+ After 2 generations at ο)=0.72. we can still see the component power-law distributions iu he Monte Carlo models. aud there is gross deviation from he analytic approximation.," After 2 generations at $Q=0.72$, we can still see the component power-law distributions in the Monte Carlo models, and there is gross deviation from the analytic approximation."
+ However. at only ο=LL. we can already see how the stochastic model is approaching he analytic version.," However, at only $n=4$, we can already see how the stochastic model is approaching the analytic version."
+ Iu panel ο at 9=21. the models are in close aerecment.," In panel $e$ at $n=24$, the models are in close agreement."
+ Iu coutrast to this rapid evolution. Figure Lidemonstratest asavt MDF ΠΟΥ) ," In contrast to this rapid evolution, Figure \ref{montecarlo}$ $f$ demonstrates the effect of a small $Q$."
+etlad Q-—0.12/ssiinilartothatinpancldforucl. Q=0.72.," The model for $n=24,\ Q=0.12$ is similar to that in panel $d$ for $n=4,\ Q=0.72$."
+ thoughthes πο ΚΟ ΟΕ ιο μμ μμ.," Although these are not identical, we can see that it is the product $nQ$ that characterizes the evolutionary state, as a result of equation \ref{binomstats}."
+..," Thus, While this statement may seem intuitively obvious, it is worth emphasizing, since the implications are profound."
+ Figure 2. shows data for the outer Galactic haloA MDF (Laird 1988: dasli«ot line). with [Fe/TII|] converted to ΟΠ following Pagel |989).," Figure \ref{halo} shows data for the outer Galactic halo MDF (Laird 1988; dash-dot line), with [Fe/H] converted to [O/H] following Pagel (1989)."
+ The solid histogram shows an inhomoegcucous Mont! Carlo model for ο)=6.1., The solid histogram shows an inhomogeneous Monte Carlo model for $nQ=6.4$.
+ There is remarkable similarity vetween the observations and our rudimentary model. espcolally in the qualitative shape. waving a metalaxich tail and large (~1.0 dex) dispersion.," There is remarkable similarity between the observations and our rudimentary model, especially in the qualitative shape, having a metal-rich tail and large $\sim 1.0$ dex) dispersion."
+" The data also agree well. coincidentally, with the Simple uodel (dotted line: e$. Tinsley 1980. equation 1.3) for xeseut-day. uretallicity log:41/Z.0.2."," The data also agree well, coincidentally, with the Simple model (dotted line; e.g., Tinsley 1980, equation 4.3) for present-day metallicity $\log z_1/\zsol=-0.2$."
+ Both models ake p= 0., Both models take $\mu_1=0$ .
+ However. it is essential to note thatprocesses.," However, it is essential to note that."
+ The hieh-uctallicity turnover in the Simple model results purely roni conusuniption of eas into stars., The high-metallicity turnover in the Simple model results purely from consumption of gas into stars.
+ Figure 3 shows the Galactic bulge. a relatively evolved svstei with data from Thata Gilmore (1995).," Figure \ref{bulge} shows the Galactic bulge, a relatively evolved system, with data from Ibata Gilmore (1995)."
+ A Simple model having :4=Z. requires frp~0.25 to resciuble the data (dotted Lue}., A Simple model having $z_1=\zsol$ requires $\mu_1\sim0.25$ to resemble the data (dotted line).
+ The dashed line shows a similar homogeneous model with no eas coustuption. 1.0. a Huear relation with equal umubers of stars created at all mictallicities (4= 1).," The dashed line shows a similar homogeneous model with no gas consumption, i.e., a linear relation with equal numbers of stars created at all metallicities $\mu_1=1$ )."
+ The difference between these two homogeneous models is thus eutirelv due to the depletion of gas into stars. which also applies to the Simple model for the halo (Figure 2)).," The difference between these two homogeneous models is thus entirely due to the depletion of gas into stars, which also applies to the Simple model for the halo (Figure \ref{halo}) )."
+ Iu coutrast. the liel-auetallicity tail of the Inhomogeneous halo model represents the vestige of the fl.) power-law: while the low-nctallicity turnover results your the Central Limit Theorem. as the component distributions progress toward gaussiuis (cf," In contrast, the high-metallicity tail of the inhomogeneous halo model represents the vestige of the $f(z)$ power-law; while the low-metallicity turnover results from the Central Limit Theorem, as the component distributions progress toward gaussians (cf."
+ Figure 1))., Figure \ref{montecarlo}) ).
+ This offers au alternative interpretation of tle AIDF vs. he Simple Model., This offers an alternative interpretation of the MDF vs. the Simple Model.
+ Ii addition. it reveals the importance of the ligh-auetallicity tail in relatively unevolved svstenis. or probing the parent distribution f(:).," In addition, it reveals the importance of the high-metallicity tail in relatively unevolved systems, for probing the parent distribution $f(z)$."
+ An evolved inhomegencous model like the Calactic »ilee should approach the same luit as the homogencous Simple model. since the metallicities coutributed by fC.) constitute a progressively simaller fraction of the ciureut : of the ISM (Ecdiuuuds 1975).," An evolved inhomogeneous model like the Galactic bulge should approach the same limit as the homogeneous Simple model, since the metallicities contributed by $f(z)$ constitute a progressively smaller fraction of the current $\bar{z}$ of the ISM (Edmunds 1975)."
+ Figure 3. indeed shows that all the models esseutially coimcide iu the low-1metallicity tail.," Figure \ref{bulge}
+ indeed shows that all the models essentially coincide in the low-metallicity tail."
+" Wowever. at the ligh-auectallicitv cud. the form of the inhomogeneous imodel mainly reflects σι, for the most recent contaminating generations."," However, at the high-metallicity end, the form of the inhomogeneous model mainly reflects $\sigma_{\rm b}$ for the most recent contaminating generations."
+ These become progressively narrower (equation 7]). but do not truncate like the Simple moclel.," These become progressively narrower (equation \ref{binomstats}) ), but do not truncate like the Simple model."
+ The inhomogeneous model matches the observed bulge MDF for jn=0.τιZ.: while the Simple model would require z4 to be 10 times larger for µη=0. with almost half the distribution at 0< [O/T <1.," The inhomogeneous model matches the observed bulge MDF for $\mu_1=0,\ z_1=\zsol$; while the Simple model would require $z_1$ to be 10 times larger for $\mu_1=0$, with almost half the distribution at $0<$ [O/H] $\lesssim 1$."
+ Thus. it is especially noteworthy that the inhomogeneous model The predicted shape. dispersion. and mean ; are linked. thus the simultancous agreement or these features is highly cucouraging.," Thus, it is especially noteworthy that the inhomogeneous model The predicted shape, dispersion, and mean $\bar{z}$ are linked, thus the simultaneous agreement for these features is highly encouraging."
+ It is also clear hat the supersolar \IDF is a vital discriminant between he models., It is also clear that the supersolar MDF is a vital discriminant between the models.
+ Malinie (1993) e1ipliasize the importance of reproducing not ouly the low-netallicitv tail. but also he lhighauctallicity drop-off of the MDE.," Malinie (1993) emphasize the importance of reproducing not only the low-metallicity tail, but also the high-metallicity drop-off of the MDF."
+ Figure L|. shows listoerams of the cuapirical Calactic hecdigketot d dso} ; neighborhood. with 1ο a," Figure \ref{gdwarf} shows histograms of the empirical Galactic disk G-dwarf MDF for stars in solar neighborhood, with [Fe/H] converted to [O/H] as before."
+nalyic inhoniogeneous node or n=32.my10ML.: the dashed line shows a 1uocel with a lower vield Wy=LAο Ιον is hereby 10 times more evolved for ie given mean mctallicity.," The solid curve shows the analytic inhomogeneous model for $nQ=32,\
+m_y=10\msol$ ; the dashed line shows a model with a lower yield $m_y=1\msol$, which is thereby 10 times more evolved for the given mean metallicity."
+ Finally. he dotted line shows 1e Simple model for zy=Z...," Finally, the dotted line shows the Simple model for $z_1=\zsol$."
+ All τιodels take jjj=0.25., All models take $\mu_1 = 0.25$.
+ The Siuple model deviates niost stronely from the data. illustrating the “Cechvart Prodem.”," The Simple model deviates most strongly from the data, illustrating the “G-dwarf Problem.”"
+ Although the ithomogereous mnodels do not solve| the Problem. their AIDEs arc| shelthy ore peaked. fms more simular to the data.," Although the inhomogeneous models do not solve the Problem, their MDFs are slightly more peaked, thus more similar to the data."
+ Note also that Rocha-Pinto Alaciel (1996) truncated their derived distribution below an effective [O/II] <0.6., Note also that Rocha-Pinto Maciel (1996) truncated their derived distribution below an effective [O/H] $< -0.6$.
+ The predicted |O/II| PWOAL dispersions in the present-day. justantancous MDE are 0.27 and 0.08 ex for the models with »Q=32 aud 320. respectively.," The predicted [O/H] FWHM dispersions in the present-day, instantaneous MDF are 0.27 and 0.08 dex for the models with $nQ = 32$ and 320, respectively."
+ These values cau be compared to the preseut-day observed ispersion iu the solar neighborhood evidenced. for example. by thewell-known 0.5 dex deviation of the Orion uchula from solar iuetallicity.," These values can be compared to the present-day observed dispersion in the solar neighborhood evidenced, for example, by thewell-known –0.3 dex deviation of the Orion nebula from solar metallicity."
+photometric error bars reported by reduction algorithms such as DoPhot (?) tend to underestimate the true photometric uncertainty for microlensing observations. which are plagued by varying observing conditions during the course of an event.,"photometric error bars reported by reduction algorithms such as DoPhot \citep{DoPhot} tend to underestimate the true photometric uncertainty for microlensing observations, which are plagued by varying observing conditions during the course of an event."
+ Apart from the systematic errorthat is taken into account for the simulated data set. underestimates of 10-20 occur frequently (e.g.222). so that the ratio X7fc.o.f. is not unlikely to exceed [4.," Apart from the systematic error that is taken into account for the simulated data set, underestimates of 10–20 occur frequently \citep[e.g.][]{OGLEerr,PLANET:first,Tsapras}, so that the ratio $\chi^2/\mathrm{d.o.f.}$ is not unlikely to exceed 1.4."
+ This strongly limits the significance of the X7 test as a measure of the goodness-of-fit., This strongly limits the significance of the $\chi^2$ test as a measure of the goodness-of-fit.
+ If one accepts an underestimate by 20'&. the single-lens model for the configurationI binarv-lens data with α=0.15 and @=0.5'A becomes acceptable. while for configuration II. single-lens models look acceptable for d.=0.15 and all applied systematic errors as well as for d=0.2 anda=1A.," If one accepts an underestimate by $20~\%$, the single-lens model for the configurationI binary-lens data with $d=0.15$ and $\sigma = 0.5\,\%$ becomes acceptable, while for configuration II, single-lens models look acceptable for $d=0.15$ and all applied systematic errors as well as for $d=0.2$ and $\sigma = 1\,\%$."
+ In addition to fits making use of the complete set of data points. best-matching single-lens models for the peak region of the lighteurve. defined as 3dx/3d. have been obtained for ay=0.3A. for which the model parameters and the result of 47 tests are displayed in Table 6..," In addition to fits making use of the complete set of data points, best-matching single-lens models for the peak region of the lightcurve, defined as $-3~\mbox{d} \leq t \leq 3~\mbox{d}$, have been obtained for $\sigma_0 = 0.3~\,\%$, for which the model parameters and the result of $\chi^2$ tests are displayed in Table \ref{tab:peak}."
+ In contrast to the fits recognizing the full data set. the baseline {ως has been fixed to its “true” value. while the blending parameter g has been allowed to vary.," In contrast to the fits recognizing the full data set, the baseline $I_\rmn{base}$ has been fixed to its 'true' value, while the blending parameter $g$ has been allowed to vary."
+ As one could have expected from the larger discrepancies being attributed to the wing regions of the light curves for configuration I. the model parameters for the fits restricted to the peak region deviate more strongly from those obtained when the full data set is considered.," As one could have expected from the larger discrepancies being attributed to the wing regions of the light curves for configuration I, the model parameters for the fits restricted to the peak region deviate more strongly from those obtained when the full data set is considered."
+ In order to adjust to the optimal model. the blending parameter g has assumed a significant non-zero value.," In order to adjust to the optimal model, the blending parameter $g$ has assumed a significant non-zero value."
+ For all selected binary-lens separations d for configuration IT and for d=0.1 for configuration I. the goodness-of-fit resulting from the 4? test for the peak region model and data is worse than for the models and data for the full light curve. so that a rejection of all considered single-lens models is indicated.," For all selected binary-lens separations $d$ for configuration II and for $d=0.1$ for configuration I, the goodness-of-fit resulting from the $\chi^2$ test for the peak region model and data is worse than for the models and data for the full light curve, so that a rejection of all considered single-lens models is indicated."
+ However. if one accounts for a possible 20% increase in the size of the photometric errors. the models with ¢@=0.1 for both configurations still survive.," However, if one accounts for a possible $20~\%$ increase in the size of the photometric errors, the models with $d=0.1$ for both configurations still survive."
+ In any case. the d.=0.1 models are not recommended for rejection for the larger systematic errors a=05% or eu=1% by means of a 47 test over the peak region.," In any case, the $d=0.1$ models are not recommended for rejection for the larger systematic errors $\sigma_0 = 0.5~\%$ or $\sigma_0 = 1\,\%$ by means of a $\chi^2$ test over the peak region."
+ Figure 6 shows a comparison of the differences between the binary-lens and the single-lens light curves of fits involving the full data set and those restricted to the peak region for a binary-lens separation parameter d.= 0.1., Figure \ref{fig:peakdel} shows a comparison of the differences between the binary-lens and the single-lens light curves of fits involving the full data set and those restricted to the peak region for a binary-lens separation parameter $d=0.1$ .
+ While the differences between these two types of fits are negligible for configuration II. the apparent differences for configuration I demonstrate the amount," While the differences between these two types of fits are negligible for configuration II, the apparent differences for configuration I demonstrate the amount"
+that it is due to the invisible companion of the eiaut detected in the crosscorrelation: even if this were the case. we would uot understand why this companion would be chromospherically active.,"that it is due to the invisible companion of the giant detected in the crosscorrelation; even if this were the case, we would not understand why this companion would be chromospherically active."
+ We conclude. that we do not uuderstaud why these two stars are N-ray sources., We conclude that we do not understand why these two stars are X-ray sources.
+ We find no indication for a faint secondary in the crosscorrelation profile of 11237: at 1e time of observation. the spectra of two cqually massive stars as suggested by Janes SSiuith (1981) woul be separated by 1 to T lan t (derived roni the ephieimieris in Mathieu et al.," We find no indication for a faint secondary in the crosscorrelation profile of 1237; at the time of observation, the spectra of two equally massive stars as suggested by Janes Smith (1984) would be separated by 4 to 7 km $^{-1}$ (derived from the ephemeris in Mathieu et al."
+ 1990)., 1990).
+ Since the secondary is 1.6 magnitude fünter iu V than the primary. we think that this small separation is compatible with finding a sinele peak iu the crosscorrelation.," Since the secondary is 1.6 magnitude fainter in V than the primary, we think that this small separation is compatible with finding a single peak in the crosscorrelation."
+ $11212 is chromospherically active. as shown by its TAIN cunission.," 1242 is chromospherically active, as shown by its K emission."
+ We sugeest that this activity. which also explains the N-ravs. is due to rapid rotation mduced by tidal interaction at periastron. which tries to bring the subeiant iuto corotation with the orbit at periastrou.," We suggest that this activity, which also explains the X-rays, is due to rapid rotation induced by tidal interaction at periastron, which tries to bring the subgiant into corotation with the orbit at periastron."
+ If we assume that the observed period of photometric variability ix the rotation period we derive au inchnation of ~ 9 Using our naxumumun value of esin; and the estinated radius., If we assume that the observed period of photometric variability is the rotation period we derive an inclination of $\sim$ $\degr$ using our maximum value of $v\sin i$ and the estimated radius.
+ This would be in agreement with a companion at the high cud of the range 0.110.91 AL. allowed by the mass fiction (Mathieu et al., This would be in agreement with a companion at the high end of the range 0.14–0.94 $M_{\odot}$ allowed by the mass function (Mathieu et al.
+ 1990)., 1990).
+ Our detection of clear chromospheric cuission indicates that the N-ray enüssion of 11010 is due to the eiaut., Our detection of clear chromospheric emission indicates that the X-ray emission of 1040 is due to the giant.
+ The white dwuf has a low temperature. aud is uulikelv to coutribute to the N-rav flux.," The white dwarf has a low temperature, and is unlikely to contribute to the X-ray flux."
+ We find a rather slow rotational velocity for the eiut. 3klkni +.," We find a rather slow rotational velocity for the giant, km $^{-1}$."
+ Ciüllilaud et (1991) detected a periodicity of 7.97 davs in the visual flux (D and V baudpasses) of S11010. with au auplitude of nunae.," Gilliland et (1991) detected a periodicity of 7.97 days in the visual flux (B and V bandpasses) of 1040, with an amplitude of mag."
+ If lis is the rotation period of the eiaut. the radius of ; (Laudsmiau ct 11997) iuplies an equatorial rotation velocity of ¢=32 kkiu 1," If this is the rotation period of the giant, the radius of $R_{\sun}$ (Landsman et 1997) implies an equatorial rotation velocity of $v=32$ km $^{-1}$."
+ This is compatible with the velocity measured with our crosscorrelatiou. esin. for an inclination /x55.37.," This is compatible with the velocity measured with our crosscorrelation, $v\sin i$, for an inclination $i\ltap5.3\degr$."
+ This inclination has ana priori probability less than:: aud it ΠΙΟΣ an unacceptablv high mass for the white cwart. from the measured mass fiction fli)=0.00268.," This inclination has an a priori probability less than; and it implies an unacceptably high mass for the white dwarf, from the measured mass function $f(m)=0.00268$."
+ We conclude that the Sddavs period cannot be the rotation period of the eiant., We conclude that the days period cannot be the rotation period of the giant.
+ It is doubtful that the white dwarf cau be responsible. as its contribution to the B aud V. flux is small.," It is doubtful that the white dwarf can be responsible, as its contribution to the $B$ and $V$ flux is small."
+ The Tea absorption profile of the blue strageler 11082 is variable., The $\alpha$ absorption profile of the blue straggler 1082 is variable.
+ If we consider the most sviuncetrie spectrum profile. that of 00:01. UT. as the unperturbed profile of the priuary. we find that the changes are duc to extra enission.," If we consider the most symmetric spectrum profile, that of 00:01 UT, as the unperturbed profile of the primary, we find that the changes are due to extra emission."
+ This is illustrated in Fie. SN., This is illustrated in Fig. \ref{s1082ha}.
+ We sueeestee that this variation is due to the subluniuous conpanion. possibly to a wh of that star.," We suggest that this variation is due to the subluminous companion, possibly to a wind of that star."
+ We have also investigated the presence of a broad shallow depression underlying the D lines (near A 5895 ÀJ) aud the triplet (near A 7775 Aj) as found by Mathys (1991)., We have also investigated the presence of a broad shallow depression underlying the D lines (near $\lambda$ 5895 ) and the triplet (near $\lambda$ 7775 ) as found by Mathys (1991).
+ We flud that tlis broad component is variable. as illustrated in Fig. 9..," We find that this broad component is variable, as illustrated in Fig. \ref{s1082na}."
+ Matlivs (1991) suggestso that the broad component originates iu the sublunünous companion., Mathys (1991) suggests that the broad component originates in the subluminous companion.
+ This companion outshines the primary by a factor six at A 1500 aud thus is prestunably hot (Landsman et al., This companion outshines the primary by a factor six at $\lambda$ 1500 and thus is presumably hot (Landsman et al.
+ 1998)., 1998).
+ We note that he star cannot be too hot or it would not show neutra lines., We note that the star cannot be too hot or it would not show neutral lines.
+While Figure Gaa shows a substantial variation in the total clistaneed traveled by the planets in these simulations but no strong dependence on the planetesimal mass.,While Figure \ref{fig:vmass}a a shows a substantial variation in the total distanced traveled by the planets in these simulations but no strong dependence on the planetesimal mass.
+ Figure Gbb. showing the semi-major axis profiles shows that the migration of the outer planet is far more stochastic when the planetesimal mass is higher as explored by previous studies (απ&Malhotra1999:Zhouetal.2002:Miurav-ClayChiang 2006).," Figure \ref{fig:vmass}b b, showing the semi-major axis profiles shows that the migration of the outer planet is far more stochastic when the planetesimal mass is higher as explored by previous studies \citep{hahn99,zhou02,murray06}."
+. In this paper we have described an implementation of the heliocentric democratic 2nd. order symplectic integration scheme (Duncanetal.1998). written in CUDA and designed to work on a graphics processing unit.," In this paper we have described an implementation of the heliocentric democratic 2nd order symplectic integration scheme \citep{duncan98}
+ written in CUDA and designed to work on a graphics processing unit."
+ This allows us to clliciently use the Large number of processors available on the card in. parallel, This allows us to efficiently use the large number of processors available on the card in parallel.
+ Phe code. is designed. to. compute a laree number of gravitational interactions as well as compute Ixeplerian advances in parallel., The code is designed to compute a large number of gravitational interactions as well as compute Keplerian advances in parallel.
+ Because all memory resides on the device this parallel computation. platform has advantages over message passage interfaces that must pav a penalty to share information between processors., Because all memory resides on the device this parallel computation platform has advantages over message passage interfaces that must pay a penalty to share information between processors.
+ The biggest drawback with our current computational platform is the low precision. a problem that is now being resolved with the recent availability of double precision capable video cards.," The biggest drawback with our current computational platform is the low precision, a problem that is now being resolved with the recent availability of double precision capable video cards."
+ We have used the new code to simulate planetary mieration into a planetesimal disk with 10 bodies and have computed. all gravitational interactions., We have used the new code to simulate planetary migration into a planetesimal disk with $10^4$ bodies and have computed all gravitational interactions.
+ We have explored, We have explored
+"at a fitted cnerey £7—6.398+0.023 kkeV having a width e=0.015(Uis kkeV and fux p=190+0.26«410? photons ?«|, (","at a fitted energy $E=6.398\pm0.023$ keV having a width $\sigma=0.045^{+0.037}_{-0.045}$ keV and flux $n=1.90 \pm 0.26 \times
+10^{-5}$ photons $^{-2}\, {\rm s}^{-1}$. ("
+Throughouto this paper. errors are quoted a confidence for the appropriate nunber of interesting paraiucters in the fit.),"Throughout this paper, errors are quoted at confidence for the appropriate number of interesting parameters in the fit.)"
+ The resulting f statistic.H was A72=130/153d.o.f.," The resulting fit statistic was $\chi^2=130/153\, d.o.f$."
+ Tutricuinely. tlie normalization of the reflection cout wis extremely hieh relative to hat of 1 primary cout.," Intriguingly, the normalization of the reflection continuum is extremely high relative to that of the primary continuum."
+ The Inconsistency of reflection and contimmiun strengths may mdicate ji the partial-covering absorption comprises a larger part of is component than currently uodeled., The inconsistency of reflection and continuum strengths may indicate that the partial-covering absorption comprises a larger part of this component than currently modeled.
+ The reatively hieh fux of the hard component was also uoted by 7) based upon fittine the 2005 d:ita. and is fully explored by ?)..," The relatively high flux of the hard component was also noted by \citet{terashima09a}
+ based upon fitting the 2005 data, and is fully explored by \citet{lobban10a}."
+ Iu adcition to t16 prominucut narrow Wy line. a strong excess of line emission is evideut im the PCA offset spectriun. at 5.Li keV (Figure 1)).," In addition to the prominent narrow K line, a strong excess of line emission is evident in the PCA offset spectrum, at 5.44 keV (Figure \ref{fig:pca}) )."
+ No sienificaut residuas appear in cigcuvector lat hat enerev (2).., No significant residuals appear in eigenvector 1 at that energy \citep{miller09b}.
+ The iw liue is thus consistent with having an origin associated. with the hard offset conrponeut and jieutral component of ]ν]να enudssion., The new line is thus consistent with having an origin associated with the hard offset component and neutral component of $\alpha$ emission.
+" As the 2008 Nov 6 epoch provides contemporaneous IIEC data for theοσα observation. we fit those spectra together over 3-SkkeV. combining the superior energysolution of TEC aud the good statistical quality of the FI NIS units (0.3) to obtain the best possible constraints on the width and enerev ofthe neutral compoucut of ίνα,"," As the 2008 Nov 6 epoch provides contemporaneous HEG data for the observation, we fit those spectra together over keV, combining the superior energy-resolution of HEG and the good statistical quality of the FI XIS units (0,3) to obtain the best possible constraints on the width and energy of the neutral component of $\alpha$."
+ A joiut fit to those data using a simple powerlaw pluρα suele Gaussian Lue showed residual excesses iu the 5 - 6 keV baud iu NIS data (Figure 2))., A joint fit to those data using a simple powerlaw plus single Gaussian line showed residual excesses in the 5 - 6 keV band in XIS data (Figure \ref{fig:cts_res2008}) ).
+ Werefit. adding to the model a single component of absorption. a Gaussian representation of the Ix.) line fixed at a vest-cncerev of kkeV with a flux linked to ]ο of the Ko. component (e.g. ?7).. a narrow absorption Lue detected at an observed enerev of T.lkkeV (as discovered by ? and confirmed bv ?)) dus a component to model the narrow line Cluission evident at 6.62kkeV iu WEC data (?)..," We refit, adding to the model a single component of absorption, a Gaussian representation of the $\beta$ line fixed at a rest-energy of keV with a flux linked to be of the $\alpha$ component \citep[e.g.][]{leahy93a,palmeri03a}, , a narrow absorption line detected at an observed energy of keV (as discovered by \citealt{pounds04c} and confirmed by \citealt{terashima09a}) ) plus a component to model the narrow line emission evident at keV in HEG data \citep{lobban10a}. ."
+" The fit vielded a line cucrey £=(6.110+41015 kkeV. width 0-26 yr ⋅⋅ ⋅↱≻↸⇅⋮⋝∶↨↸∖∖⋜⋯≼↧∐∪↥⋅↕⊔⋜↧∐∑⋜↧⊓∪∐∣∣∶↓∙∩∩∶∶↽ν↽10""photounscenmD5 Lo "," The fit yielded a line energy $E=6.410 \pm 0.015$ keV, width $\sigma=50^{+26}_{-33}$ eV and normalization $n=1.90\pm 0.395 \times 10^{-5} {\rm photons\, cm^{-2}\, s^{-1}}$ ."
+Τ]μο equivadeut width for the FelkI&a liue componcw is 9heeV agaist he total contimmiun level observed during Ww08., The equivalent width for the $\alpha$ line component is eV against the total continuum level observed during 2008.
+" The constraint ou the width of the Fe Ίνα line is equivalent to à FWIIM velocity-broacdening of 5510P | (00ppovento conufideuce interval),", The constraint on the width of the Fe $\alpha$ line is equivalent to a FWHM velocity-broadening of $5540_{-3664}^{+2846}$ $^{-1}$ percent confidence interval).
+" Asstming the true lnc| ceutroid enerev to be kkeV. the coufidence constraint on euergv from the combined TEC""NIS data limits the bulls velocity. of hisfo cutterqul- to or250Ne¢1200 + where negative velocity denotes outflow."," Assuming the true line centroid energy to be keV, the confidence constraint on energy from the combined HEG/XIS data limits the bulk velocity of this emitter to $250 \ga v \ga -1200$ $^{-1}$ where negative velocity denotes outflow."
+ As the PCA decomposition iudicates the μμ... of a weak broad component of Fe Ίνα on eigenvector ] (7). we vied adding a second componeut of Fe Ika eudssiou to the model. with the euergies of both liue components lukeclL.," As the PCA decomposition indicates the presence of a weak broad component of Fe $\alpha$ on eigenvector 1 \citep{miller09b}, we tried adding a second component of Fe $\alpha$ emission to the model, with the energies of both line components linked."
+" This acldition inuproved the fit-statistic Jw AV5=5.Γ givingD. the same line euergv as for he single component model. with a,=Ebod eV and normalization my=1.1240.32410.""photonscu17s ηll eV) oy=Lily eV and normalization no=1.070.15«10.°photouscm7s I(EWaBey)."," This addition improved the fit-statistic by $\Delta \chi^2=5$, giving the same line energy as for the single component model, with $\sigma_1 =1^{+23}_{-1}$ eV and normalization $n_1=1.12\pm 0.32 \times 10^{-5} {\rm photons\, cm^{-2}\, s^{-1}}$ $=41$ eV), $\sigma_2 =145^{+70}_{-52}$ eV and normalization $n_2=1.07\pm 0.45 \times 10^{-5} {\rm photons\, cm^{-2}\, s^{-1}}$ $=43$ eV)."
+ Although he two-component parameterization of FekKlIkao offers only a niuginal iniprovement to the fit. it is of interest with regard to consisteucv with the PCA decomposition. which indicates the weak. broad component of Fe Io eumisslon to be present on eieenvector 1.," Although the two-component parameterization of $\alpha$ offers only a marginal improvement to the fit, it is of interest with regard to consistency with the PCA decomposition, which indicates the weak broad component of Fe $\alpha$ emission to be present on eigenvector 1."
+ Iu subsequeu fits we fix the Fe ka inodel liue width at σ=50 eV obained from the simple fit. while carne in muiud he possible more complex solution for the line profile.," In subsequent fits we fix the Fe $\alpha$ model line width at $\sigma=50$ eV obtained from the simple fit, while bearing in mind the possible more complex solution for the line profile."
+" As the PCA is consistent with a common origim for Fe Ikea Cluission and the line at 5.11 τον, and as the widths of the two lines are cousistent (aud as there is no other useful constrain available for the indeterminate line width) hereater we fixed both tog=50 eV in spectral fitting."," As the PCA is consistent with a common origin for Fe $\alpha$ emission and the line at 5.44 keV, and as the widths of the two lines are consistent (and as there is no other useful constraint available for the indeterminate line width) hereafter we fixed both to $\sigma=50$ eV in spectral fitting."
+ To assess the sieuificance of the line at 5.11 keV owe returned to direct fittine of the 2005 data. where the source was observed ο be ata low flux aud where PCA iudicates that the line would have the highest equivaleut width.," To assess the significance of the line at 5.44 keV we returned to direct fitting of the 2005 data, where the source was observed to be at a low flux and where PCA indicates that the line would have the highest equivalent width."
+ We fit the suuunaued 2005 data frou the FI NIS units 0. 2and 3 using an absorbed οςutiuuuuu model. applving it to the 3-LOkkeV bauL for the purpose of examining the data for featurOs of interest.," We fit the summed 2005 data from the FI XIS units 0, 2and 3 using an absorbed continuum model, applying it to the keV band for the purpose of examining the data for features of interest."
+ ACaussian line was included in he fit. to account," AGaussian line was included in the fit, to account"
+cuuission in the region associated to the WALIAP haze.,emission in the region associated to the WMAP haze.
+ A fullv sel£cousisteut galactic cosmic-ray model Is necessary fo model the astroplivsical diffuse ciission from the Galaxy iud to compare it to the Fermi-LAT data., A fully self-consistent galactic cosmic-ray model is necessary to model the astrophysical diffuse emission from the Galaxy and to compare it to the Fermi-LAT data.
+" We thank Trov Porter and Andy Strong for feedback regarding models. as well as Cregory Dobler and Douglas Finkbeimer for useful discussions regarding the template construction in DOO,"," We thank Troy Porter and Andy Strong for feedback regarding models, as well as Gregory Dobler and Douglas Finkbeiner for useful discussions regarding the template construction in D09."
+ TL is supported by a GAANN Fellowship by the Departineut of Education., TL is supported by a GAANN Fellowship by the Department of Education.
+ SP is supported by an ΟΠ Award from the US Departineut of Eucrgv (DoE Coutract DEFCO2-0LER11268). aud by NSF Craut PITY-0757911.," SP is supported by an OJI Award from the US Department of Energy (DoE Contract DEFG02-04ER41268), and by NSF Grant PHY-0757911."
+marking (heir trajectories in (he absence of deflections.,marking their trajectories in the absence of deflections.
+ The peak in the density of intersecting trajectories provides the best estimator for the point of origin of the measured motions., The peak in the density of intersecting trajectories provides the best estimator for the point of origin of the measured motions.
+ The “intensity of each (rail is weighted by the speed., The “intensity” of each trail is weighted by the speed.
+ Thus. large motions. which can be measured more precisely than small motions. have a greater weight.," Thus, large motions, which can be measured more precisely than small motions, have a greater weight."
+" The image in the rieht panel has been boxcar-smoothed by 31 pixels (25"") to aid visualization of the intersection point.", The image in the right panel has been boxcar-smoothed by 31 pixels ) to aid visualization of the intersection point.
+" The point of origin of the proper motions is located at J(2000) = 5 35"" 1455. 23"". with alo uncertainty of about5""."," The point of origin of the proper motions is located at J(2000) = $^h$ $^m$ $^s$ 5, with a 1 $\sigma$ uncertainty of about."
+ This location is offset by about to the northeast of the apparent ejection center of the massive stars and nnorth of the eurrent location of radio source I. Zapata et al. (, This location is offset by about to the northeast of the apparent ejection center of the massive stars and north of the current location of radio source I. Zapata et al. (
+"2010) found an expansion center based on interferometric CO J = 21 observations at J(2000) = 5 35"" 14537. 27.""9, about delve south the expansion center derived from our proper motions. but consistent within the uncer(aintiv estimates.","2010) found an expansion center based on interferometric CO J = 2–1 observations at J(2000) = $^h$ $^m$ $^s$ 37, 9, about due south the expansion center derived from our proper motions, but consistent within the uncertainty estimates."
+ lligh velocity runaway stars are relatively common among massive O and early D stars but rare among lower-nmmss stars (Giles Bolton 1986: Gies 1987)., High velocity runaway stars are relatively common among massive O and early B stars but rare among lower-mass stars (Gies Bolton 1986; Gies 1987).
+ While approximately of O stars are runaways with velocities ereater than 30 km s.| with respect to their birth clusters ancl associations. less than of Band A stars have such velocities (Stone 1991: Sehilbach Roser 2008: Zinnecker Yorke 2007).," While approximately of O stars are runaways with velocities greater than 30 km $^{-1}$ with respect to their birth clusters and associations, less than of B and A stars have such velocities (Stone 1991; Schilbach Roser 2008; Zinnecker Yorke 2007)."
+ Many. (possibly most) massive runaway stars were ejected bv three-bodsy. encounters between a single star and a massive binary (Gvaramadze Gualandris 2010). lour-body encounters between a pair of massive binaries (Poveda. Ruiz. Allen 1967). or the decay of non-hierarchical multiple star svstems (van Albada 1968: Valtonen Mikkola 1991: Aarseth et al.," Many (possibly most) massive runaway stars were ejected by three-body encounters between a single star and a massive binary (Gvaramadze Gualandris 2010), four-body encounters between a pair of massive binaries (Poveda, Ruiz, Allen 1967), or the decay of non-hierarchical multiple star systems (van Albada 1968; Valtonen Mikkola 1991; Aarseth et al."
+ 1994: Sterzik Durisen 1995. 1993).," 1994; Sterzik Durisen 1995, 1998)."
+ Such three- or four-bodsy encounters are most likely (o occur in very dense. voung star clusters.," Such three- or four-body encounters are most likely to occur in very dense, young star clusters."
+" The runaway stars AE Aur and ji Col. ejected with speeds of over 100 kms + about 2.5xLO"" vears ago (Lloogerwerl. de Bruijne. de Zeeuw 2001) from the NGC 1980 cluster located about south of the Orion Nebula. provide a prototypical example."," The runaway stars AE Aur and $\mu$ Col, ejected with speeds of over 100 km $^{-1}$ about $2.5 \times 10^6$ years ago (Hoogerwerf, de Bruijne, de Zeeuw 2001) from the NGC 1980 cluster located about south of the Orion Nebula, provide a prototypical example."
+ This event produced the massive colliding-wind X-ray binary + Ori (Gualanclris. Portegies-Zwart. Egeleton 2004).," This event produced the massive colliding-wind X-ray binary $\iota$ Ori (Gualandris, Portegies-Zwart, Eggleton 2004)."
+ Non-hierarchical multiple star svstems. triple or hieher-order multiples in which the lime-averaged clistances between components are similar are dynamically unstable.," Non-hierarchical multiple star systems, triple or higher-order multiples in which the time-averaged distances between components are similar are dynamically unstable."
+ Numerical experiments show that non-hierarchical svstems tend to rearrange into hierarchical configurations, Numerical experiments show that non-hierarchical systems tend to rearrange into hierarchical configurations
+Note that is integer for the solutions. which vanish al the poles (4¢=41).,"Note that $n$ is integer for the solutions, which vanish at the poles $\mu=\pm 1$ )."
+ The dispersion relation is found to be The dispersion relation can be split into high (magnetic Poincaré)) aud low (magnetic Rossby) frequency branches., The dispersion relation is found to be The dispersion relation can be split into high (magnetic Poincaré)) and low (magnetic Rossby) frequency branches.
+ For the magnetic Poincaré waves we have This is the dispersion relation of ordinary surlace eravily waves and we can recover il for the high s harmonics (Longuet-Iligeins1965)., For the magnetic Poincaré waves we have This is the dispersion relation of ordinary surface gravity waves and we can recover it for the high $s$ harmonics \citep{lon65}.
+. For the magnetic Rossby waves we have approximately It is clear from equation (6) that the magnetic field causes the splitting of ordinary Rosshv waves into [ast and slow modes (Zaqaras, For the magnetic Rossby waves we have approximately It is clear from equation (6) that the magnetic field causes the splitting of ordinary Rossby waves into fast and slow modes \citep{zaq07}.
+hvilietal.2007).. For ay«1 the dispersion relations for the fast and slow magnetic Rossby modes can be approximated as We see (hal the last magnetic Rossby mode has a dispersion relation similar to that ol the ordinary Rossby waves but slightly modified by the magnetic field., For $\alpha_0 \ll 1$ the dispersion relations for the fast and slow magnetic Rossby modes can be approximated as We see that the fast magnetic Rossby mode has a dispersion relation similar to that of the ordinary Rossby waves but slightly modified by the magnetic field.
+ Slow magnetic Rossby waves. however. are new wave modes as (heir dispersion relation is different fom Rossby and Alfvénn wave dispersion relations (the dispersion relation of Alfvénn waves should be Ay= say).," Slow magnetic Rossby waves, however, are new wave modes as their dispersion relation is different from Rossby and Alfvénn wave dispersion relations (the dispersion relation of Alfvénn waves should be $\lambda_A={\pm}s\alpha_0$ )."
+ There is a significant difference between the frequencies of fast and slow magnetic Rossby modes for sufficiently small ay., There is a significant difference between the frequencies of fast and slow magnetic Rossby modes for sufficiently small $\alpha_0$.
+ The dispersion diagram for the s=1 harmonics of shallow water MIID waves accordingto Eq. (41) , The dispersion diagram for the $s=1$ harmonics of shallow water MHD waves accordingto Eq. \ref{dr_small_eps}) )
+is shown in Fie 1., is shown in Fig 1.
+ Here e=0.01 and ay= 0.05., Here $\epsilon = 0.01$ and $\alpha_0=0.05$ .
+ The frequencies of magnetic, The frequencies of magnetic
+Results froma variety of solar. aümospheric and terrestrial neutrino oscillation experiments 4] have constrained the form of the lepton mixing matrix C |2].,"Results from a variety of solar, atmospheric and terrestrial neutrino oscillation experiments \cite{1} have constrained the form of the lepton mixing matrix $U$ \cite{2}."
+. The lepton mixing matrix is given bv, The lepton mixing matrix is given by
+et al.,et al.
+ 2004. Isieh et al.," 2004, Hsieh et al."
+ 2010). qualitatively consistent with sublimation in the same wav (hat mass loss from the classical comets is modulated by the varving insolation around the orbit.," 2010), qualitatively consistent with sublimation in the same way that mass loss from the classical comets is modulated by the varying insolation around the orbit."
+ Other mechanisms in Section (3)) are either inconsistent with recurrent mass loss (impact. rotational instabilitv). or ean be reconciled with it only in the most contrived wav.," Other mechanisms in Section \ref{mechanisms}) ) are either inconsistent with recurrent mass loss (impact, rotational instability), or can be reconciled with it only in the most contrived way."
+ With equilibrium sublimation rates <(3 to x10. > kem 7s bat R= 2.6 to 3.0 AU (Figure 4)). the most stringent. production rates <O0.01 kg ! (Table 2)) imply mass loss [rom areas of exposed ice <200 to 300 ni. corresponding to circles of radius ~10 m (IIsieh et al.," With equilibrium sublimation rates $\lesssim$ (3 to $\times$ $^{-5}$ kg $^{-2}$ $^{-1}$ at $R$ = 2.6 to 3.0 AU (Figure \ref{sublimation}) ), the most stringent production rates $\lesssim$ 0.01 kg $^{-1}$ (Table \ref{physical}) ) imply mass loss from areas of exposed ice $\lesssim$ 200 to 300 $^{2}$, corresponding to circles of radius $\sim$ 10 m (Hsieh et al."
+ 2004)., 2004).
+ The eccentricities of 133P. and 238P are smiiadl (0.165. ancl 0.253. respectively) and. over the range of heliocentric distances swept. the expected sublimation mass flux varies by less than an order of magnitude (upper curve in Figure 4)).," The eccentricities of 133P and 238P are small (0.165 and 0.253, respectively) and, over the range of heliocentric distances swept, the expected sublimation mass flux varies by less than an order of magnitude (upper curve in Figure \ref{sublimation}) )."
+ This is a problem because a ten-fold reduction in activity al aplielion relative to perihelion would still produce detectable dust. whereas none is observed.," This is a problem because a ten-fold reduction in activity at aphelion relative to perihelion would still produce detectable dust, whereas none is observed."
+ Instead. mass loss might be confined to a window near perihelion because of a shadowing elfect of local topography on the nucleus. perhaps caused by recession of the sublimating surface. forming a pit (Equation 3)).," Instead, mass loss might be confined to a window near perihelion because of a shadowing effect of local topography on the nucleus, perhaps caused by recession of the sublimating surface, forming a pit (Equation \ref{recession}) )."
+ The surface recession rate al A = 2.6 to 3.0 AU is (3 to x  ms I. so that a sublimation pit meters deep should develop in only a lew orbits.," The surface recession rate at $R$ = 2.6 to 3.0 AU is (3 to $\times$ $^{-8}$ m $^{-1}$, so that a sublimation pit meters deep should develop in only a few orbits."
+ Llowever. that the ice must still be close to the physical surface is shown bv the appearance of coma near and soon after perihelion.," However, that the ice must still be close to the physical surface is shown by the appearance of coma near and soon after perihelion."
+ Ice more (han a lew centimeters deep would be thermally decoupled [rom the deposition of heat on the surface., Ice more than a few centimeters deep would be thermally decoupled from the deposition of heat on the surface.
+ The evidence for ice sublimation. although self-consistent. remains indirect.," The evidence for ice sublimation, although self-consistent, remains indirect."
+ Detection of easeous products of sublimation would provide definitive evidence for (he role of sublimation in active asteroids., Detection of gaseous products of sublimation would provide definitive evidence for the role of sublimation in active asteroids.
+ Just as in the classical comets. practical concerns render (he daughter species (molecular fragments) optically more visible (han the parent molecules sublimated from the ice.," Just as in the classical comets, practical concerns render the daughter species (molecular fragments) optically more visible than the parent molecules sublimated from the ice."
+ The most prominent band in (he optical spectra of comets is due to the trace species CN. itself a product of dissociation of the parent molecule ΗΝ.," The most prominent band in the optical spectra of comets is due to the trace species CN, itself a product of dissociation of the parent molecule HCN."
+ As noted in Section 2. several attempts to detect CN in active asteroids have been reported. (vpically resulting in limits to the total gas production rate 0.1 - 1 kg |. about two to three orders of magnitude smaller than found in bright comets near | AU.," As noted in Section 2, several attempts to detect CN in active asteroids have been reported, typically resulting in limits to the total gas production rate 0.1 - 1 kg $^{-1}$, about two to three orders of magnitude smaller than found in bright comets near 1 AU."
+ To set this on a quantitative basis. we show in Figure (7)) the expected CN. band Εν as a function of δὲ for several outgassing rates. as marked. computed as described in Jewitt et al. (," To set this on a quantitative basis, we show in Figure \ref{cn_plot}) ) the expected CN band flux as a function of $R$ for several outgassing rates, as marked, computed as described in Jewitt et al. ("
+2009).,2009).
+" The fluxes in the Figure were evaluated lor opposition (ie..N = H- 1) and assume the IHaser model and a spectroscopic slit 1x5"".", The fluxes in the Figure were evaluated for opposition (i.e. $\Delta$ = $R$ - 1) and assume the Haser model and a spectroscopic slit $\arcsec \times$ $\arcsec$.
+ The Swines effect (in which the resonance fluorescence cross-section is affected by the heliocentric velocity through the Doppler effect) amounts to less than a factor of two. and has been neglected [or simplicity.," The Swings effect (in which the resonance fluorescence cross-section is affected by the heliocentric velocity through the Doppler effect) amounts to less than a factor of two, and has been neglected for simplicity."
+ Plotted on the figure are published limits to the CN band flux in active asteroids., Plotted on the figure are published limits to the CN band flux in active asteroids.
+" Slate of ihe art measurements. indicated bv the horizontal grey band in the figure. reach CN fIuxes ~ (210 3)x10 ! erg 7s 1!, corresponding to Qe,~ 1075 tat R= 2 AU but. by 9 "," State of the art measurements, indicated by the horizontal grey band in the figure, reach CN fluxes $\sim$ (2 to $\times$ $^{-16}$ erg $^{-2}$ $^{-1}$, corresponding to $Q_{CN} \sim$ $^{23}$ $^{-1}$ at $R$ = 2 AU but, by $R$ "
+GHz.,GHz.
+ This scaling will be continued throughout., This scaling will be continued throughout.
+ The Doppler axis is scaled by rz/V.g., The Doppler axis is scaled by $r_\text{f} / V_\text{eff}$.
+ With the values used earlier the Doppler range would be +3 mHz., With the values used earlier the Doppler range would be $\pm3$ mHz.
+" An arc is visible in Figure 7, but it is quite symmetric and not sharply defined."," An arc is visible in Figure 7, but it is quite symmetric and not sharply defined."
+ Most observed arcs are more clearly defined and they are often quite asymmetric., Most observed arcs are more clearly defined and they are often quite asymmetric.
+ It is known that the arcs would be sharper in weaker scintillation or if the scattering were anisotropic (Cordes et al., It is known that the arcs would be sharper in weaker scintillation or if the scattering were anisotropic (Cordes et al.
+ 2006) and this is a strong argument that interstellar turbulence is often anisotropic., 2006) and this is a strong argument that interstellar turbulence is often anisotropic.
+ However it has not been understood exactly why observations often show asymmetric arcs., However it has not been understood exactly why observations often show asymmetric arcs.
+" It has been assumed that this is due, in some way, to refraction, but it has been unclear whether this requires a discrete refracting structure, or whether the refraction that occurs naturally due to the low frequency part of the turbulent spectrum is sufficient."," It has been assumed that this is due, in some way, to refraction, but it has been unclear whether this requires a discrete refracting structure, or whether the refraction that occurs naturally due to the low frequency part of the turbulent spectrum is sufficient."
+" To test this question we have broken the full dynamic spectrum into smaller blocks, typical of the observation duration (as was done in calculating the time and frequency scales in Figure 3)."," To test this question we have broken the full dynamic spectrum into smaller blocks, typical of the observation duration (as was done in calculating the time and frequency scales in Figure 3)."
+" 'The results, which are shown in Figure 8, show exactly the characteristics of the observations."," The results, which are shown in Figure 8, show exactly the characteristics of the observations."
+ The arcs are sometimes quite asymmetric and the orientation of the asymmetry reverses on a time scale considerably longer than the refractive time scale., The arcs are sometimes quite asymmetric and the orientation of the asymmetry reverses on a time scale considerably longer than the refractive time scale.
+" Evidently this resolves the question of whether such asymmetric secondary spectra are due to discrete deterministic structures or to refractive components of a turbulent spectrum - they can be caused by a continuous turbulent spectrum, but the scales involved are larger than the refractive scale."," Evidently this resolves the question of whether such asymmetric secondary spectra are due to discrete deterministic structures or to refractive components of a turbulent spectrum - they can be caused by a continuous turbulent spectrum, but the scales involved are larger than the refractive scale."
+" When the same phase screen is simulated without dispersion the dynamic spectrum, which is plotted in Figure 1(c), shows criss-cross structures but very little tilt."," When the same phase screen is simulated without dispersion the dynamic spectrum, which is plotted in Figure 1(c), shows criss-cross structures but very little tilt."
+" The average secondary spectrum is very similar to that shown in Figure 7, but the secondary spectra of shorter blocks, which are shown in Figure 9, do not show the time varying asymmetry characteristic of the dispersive simulation shown in Figure 8."," The average secondary spectrum is very similar to that shown in Figure 7, but the secondary spectra of shorter blocks, which are shown in Figure 9, do not show the time varying asymmetry characteristic of the dispersive simulation shown in Figure 8."
+" Thus there can be no doubt that the phenomenon of tilted bands in dynamic spectra is not caused solely by refraction, but by dispersion and refraction."," Thus there can be no doubt that the phenomenon of tilted bands in dynamic spectra is not caused solely by refraction, but by dispersion and refraction."
+ With this insight it is relatively simple to describe the phenomenon analytically., With this insight it is relatively simple to describe the phenomenon analytically.
+ The physical difference can be explained intuitively., The physical difference can be explained intuitively.
+ Large scale refraction will tilt an entire angular spectrum regardless of dispersion., Large scale refraction will tilt an entire angular spectrum regardless of dispersion.
+ However it will tilt to the minimum phase delay position (due to Fermat’s Principle)., However it will tilt to the minimum phase delay position (due to Fermat's Principle).
+ The secondary spectrum is caused by interference of scattered waves with different group delay and different Doppler shift., The secondary spectrum is caused by interference of scattered waves with different group delay and different Doppler shift.
+ In a non-dispersive problem the group delay is equal to the phase delay so a refractive shift of the entire angular spectrum has no effect on the secondary spectrum., In a non-dispersive problem the group delay is equal to the phase delay so a refractive shift of the entire angular spectrum has no effect on the secondary spectrum.
+" When the medium is dispersive the phase delay is the negative of the group delay, so the position of minimum phase delay is not the position of minimum group delay."," When the medium is dispersive the phase delay is the negative of the group delay, so the position of minimum phase delay is not the position of minimum group delay."
+ Indeed it is a local maximum of group delay., Indeed it is a local maximum of group delay.
+ In this case a refractive tilt has a dramatic effect on the parabolic arc., In this case a refractive tilt has a dramatic effect on the parabolic arc.
+ A rough analysis is given below., A rough analysis is given below.
+ We assume (without loss of generality) a thin screen half way between a pulsar and the observer., We assume (without loss of generality) a thin screen half way between a pulsar and the observer.
+ The arcs are caused by the interference of a wave scattered at an angle 60 in the direction of the velocity and an unscattered wave., The arcs are caused by the interference of a wave scattered at an angle $\delta\theta$ in the direction of the velocity and an unscattered wave.
+ The two waves interfere with differential Doppler shiftsδι=2V60/A and differential delays ὅτα=200?/c where z is the distance from the screen to the observer., The two waves interfere with differential Doppler shifts$\delta f_d = 2V \delta\theta/\lambda$ and differential delays $\delta \tau_d = z\delta\theta^2/c$ where z is the distance from the screen to the observer.
+" It is useful to normalize the parameters with respect to the rms scattering angle 0o, i.e. fag=2V69/A and Tao= 202/c."," It is useful to normalize the parameters with respect to the rms scattering angle $\theta_0$, i.e. $f_{d0} = 2V \theta_0/\lambda$ and $\tau_{d0} = z\theta_0^2/c$ ."
+ Then ὅτα/ταρ=(Ofa/fao)? defines the parabolic, Then $\delta \tau_d/\tau_{d0} = (\delta f_d/f_{d0})^2$ defines the parabolic
+indicate that the size of this absorber is of the order of the size of giant hydrogen clouds in our galaxy or of a galactie disk seen face-on.,indicate that the size of this absorber is of the order of the size of giant hydrogen clouds in our galaxy or of a galactic disk seen face-on.
+ If this is not a galactic disk. at these redshift. systems of this scale could be protogalactic clumps. that is the building blocks of the various type of galaxies (ellipticals. spirals. etc.)," If this is not a galactic disk, at these redshift, systems of this scale could be protogalactic clumps, that is the building blocks of the various type of galaxies (ellipticals, spirals, etc.)"
+ that are observed at present epoch (Khersonsky 1996.. Haehneltetal. 1998)).," that are observed at present epoch \cite{Khersonsky96}, \cite{Haehnelt98}) )."
+ In the damped Lya system at 2=3.0513 of 2619 studied by SOM absorption was detected., In the damped $\alpha$ system at $z=3.0543$ of $-$ 2619 studied by SOM absorption was detected.
+ In this paper we report the detection at 3.50 of absorption from the excited fine structure level of at 2=3.0513., In this paper we report the detection at $\sigma$ of absorption from the excited fine structure level of at $z=3.0543$.
+ From the measure of the equivalent width of the two lines we derived an upper limit of 21.6 K on the fine structure excitation temperature., From the measure of the equivalent width of the two lines we derived an upper limit of 21.6 K on the fine structure excitation temperature.
+ This value provides a strict upper limit on the temperature of the CMB at;=3.0513. which at this redshift is predicted to be 11.05 K. We then used the derived relative populations of the fine structure levels of to set constraints on the absorber density and on the UV field in the absorbing cloud.," This value provides a strict upper limit on the temperature of the CMB at $z=3.0543$, which at this redshift is predicted to be 11.05 K. We then used the derived relative populations of the fine structure levels of to set constraints on the absorber density and on the UV field in the absorbing cloud."
+ Assuming an 10nization degree ranging from 0 to10%.. a gas kinetic temperature between 100 and 100000 K and a UV field with a Milky Way spectrum.= the density of the absorber is constrained to be between 0.7 and 40  and the H-ionizing flux between 1«10 and διςLOScm2s+. that is between | and 80 times the intensity of the Galactic UV field.," Assuming an ionization degree ranging from 0 to, a gas kinetic temperature between 100 and 000 K and a UV field with a Milky Way spectrum, the density of the absorber is constrained to be between 0.7 and 40 $^{-3}$ and the H-ionizing flux between $1\times 10^7$ and $8\times
+10^8~{\rm cm^{-2}s^{-1}}$, that is between 1 and 80 times the intensity of the Galactic UV field."
+ The upper limits hold even if the detected absorption from I at 2=3.0513 1s seriously contaminated by Lyman forest absorption., The upper limits hold even if the detected absorption from $^*$ at $z=3.0543$ is seriously contaminated by Lyman forest absorption.
+ If the damped Lya system Is assumed to be homogeneous. a density value between 0.7 and 40 Ὁ constrains the size of this absorber at +—3 to be between | and 100 pe in the direction of the line of sight.," If the damped $\alpha$ system is assumed to be homogeneous, a density value between 0.7 and 40 $^{-3}$ constrains the size of this absorber at $z \sim 3$ to be between 1 and 100 pc in the direction of the line of sight."
+"interesting and relatively compact feature at [R.A., dec] c ","interesting and relatively compact feature at [R.A., dec] $\approx [133.9\arcdeg, 3.6\arcdeg$ ]."
+"The object contains ~300 stars to g—22, has a [133.9?,3.6?].filtered star count density higher than any visible portion of the EBS, and appears quite distinct within the stream."," The object contains $\approx 300$ stars to $g = 22$, has a filtered star count density higher than any visible portion of the EBS, and appears quite distinct within the stream."
+" There are several background galaxy clusters identified within 1? of this position, but examination of the identically filtered SDSS DR7 galaxy catalog shows no significant galaxy concentration of similar size or shape."," There are several background galaxy clusters identified within $1\arcdeg$ of this position, but examination of the identically filtered SDSS DR7 galaxy catalog shows no significant galaxy concentration of similar size or shape."
+" T'wo lesser peaks are apparent some 4? north and south of this object (see also Figure 4), but we focus on this object because it is the most prominant and populous concentration, both to the eye and in the longitudonal profile."," Two lesser peaks are apparent some $4\arcdeg$ north and south of this object (see also Figure 4), but we focus on this object because it is the most prominant and populous concentration, both to the eye and in the longitudonal profile."
+" Given its apparent location at the same distance as the stream and approximately centered within it, we infer that the feature is physically associated with the EBS and we designate it Hydra I. The feature appears somewhat amorphous, with two primary concentrations extending to the north and west, respectively, for a total extent of zz2°."," Given its apparent location at the same distance as the stream and approximately centered within it, we infer that the feature is physically associated with the EBS and we designate it Hydra I. The feature appears somewhat amorphous, with two primary concentrations extending to the north and west, respectively, for a total extent of $\approx 2\arcdeg$."
+" At 9.7 kpc this corresponds to a spatial extent of about ~350 pc, which is far larger than any known globular cluster."," At 9.7 kpc this corresponds to a spatial extent of about $\sim 350$ pc, which is far larger than any known globular cluster."
+" Using filter shifts to estimate relative distances, we find that the maximum filtered surface densities in the two lobes of Hydra I occur within 0.1 magnitudes of one another, indicating that Hydra I is not significantly extended along our line of sight."," Using filter shifts to estimate relative distances, we find that the maximum filtered surface densities in the two lobes of Hydra I occur within 0.1 magnitudes of one another, indicating that Hydra I is not significantly extended along our line of sight."
+" At 9.7 kpc, 0.1 magnitudes corresponds to a difference in distance of 500 pc."," At 9.7 kpc, 0.1 magnitudes corresponds to a difference in distance of 500 pc."
+" To within the uncertainties, this is identical to the lateral extent of the object."," To within the uncertainties, this is identical to the lateral extent of the object."
+" How real is the apparent, double-lobed morphology of Hydra I?"," How real is the apparent, double-lobed morphology of Hydra I?"
+ Could the northeastern lobe be simply a chance consequence of Poisson statistics?, Could the northeastern lobe be simply a chance consequence of Poisson statistics?
+" The surface density profile of stars with 19«g23 in the stronger, western lobe and lying within 0.2 magnitudes of the Z=0.0003 ischrone in g—4 is shown in Figure 6."," The surface density profile of stars with $19 < g < 23$ in the stronger, western lobe and lying within 0.2 magnitudes of the Z=0.0003 ischrone in $g -i$ is shown in Figure 6."
+" Fitting an elliptical 1/r model to the star counts, we find ellipticity e=0.45+0.05, with 0=90+10? (measured north through east), a total population out to 1? of 300+10 stars, and an overall x? of 1.2."," Fitting an elliptical $1/r$ model to the star counts, we find ellipticity $\epsilon = 0.45 \pm 0.05$ , with $\theta = 90 \pm 10\arcdeg$ (measured north through east), a total population out to $1\arcdeg$ of $300 \pm
+10$ stars, and an overall $\chi^2$ of 1.2."
+" Following Martin(DeJong), we then determine the fractional R.M.S. deviation o,-/total of the data compared to the model."," Following \citet{martin2008}, we then determine the fractional R.M.S. deviation $\sigma_{sc}/$ total of the data compared to the model."
+" We generate 1000 Poisson realizations of the field out to 1° and examine the distribution of (c,./total)?.", We generate 1000 Poisson realizations of the field out to $1\arcdeg$ and examine the distribution of $\sigma_{sc}$ $^2$.
+" The peak of the distribution differs from 0 at the 9c level, indicating that the northeastern lobe is unlikely to be a statistical departure from the model."," The peak of the distribution differs from 0 at the $9\sigma$ level, indicating that the northeastern lobe is unlikely to be a statistical departure from the model."
+ We can only speculate as to the relationship between the two lobes at this point., We can only speculate as to the relationship between the two lobes at this point.
+" However, being part of a tidal stream, it seems reasonable to suppose that the two lobes are unbound, comoving tidal remnants."," However, being part of a tidal stream, it seems reasonable to suppose that the two lobes are unbound, comoving tidal remnants."
+ Deeper imaging of Hydra I is currently being acquired and will be the subject of a future contribution., Deeper imaging of Hydra I is currently being acquired and will be the subject of a future contribution.
+" Schlaufmanetal. detected a velocity overdensity among metal (2009)poor |stars at [R.A., dec] — [134°, 3.2°] that they attributed to the Monoceros ring."," \citet{schlaufman2009} detected a velocity overdensity among metal poor stars at [R.A., dec] = $134\arcdeg, 3.2\arcdeg$ ] that they attributed to the Monoceros ring."
+" Figure 5 shows that the stars making up their B-8/PCI-9/PCII-21 detection clearly sample the position of Hydra I, as well as à portion of the EBS extending to the south (which would be expected to have a nearly identical velocity)."," Figure 5 shows that the stars making up their B-8/PCI-9/PCII-21 detection clearly sample the position of Hydra I, as well as a portion of the EBS extending to the south (which would be expected to have a nearly identical velocity)."
+ Their estimated distance of 9.7 kpc is identical with the values we find above., Their estimated distance of 9.7 kpc is identical with the values we find above.
+ Filtering this region of DR7 using isochrones with metallicities spanning the range —2.2« [Fe/H]«0.0 reveals no other significant structures at this position., Filtering this region of DR7 using isochrones with metallicities spanning the range $-2.2 < $ $ < 0.0$ reveals no other significant structures at this position.
+ We conclude that Schlaufmanetal.(2009) actually sampled the stars in Hydra I and the EBS., We conclude that \citet{schlaufman2009} actually sampled the stars in Hydra I and the EBS.
+ They find à mean velocity for B-8/PCI-9/PCII-21 of 85 km s! and a dispersion of 14.9 km s~!., They find a mean velocity for B-8/PCI-9/PCII-21 of 85 km $^{-1}$ and a dispersion of 14.9 km $^{-1}$.
+" The velocity dispersion is again quite high compared with globular clusters or ultrafaint galaxies, suggesting that Hydra I was either quite massive or has been significantly heated over time."," The velocity dispersion is again quite high compared with globular clusters or ultrafaint galaxies, suggesting that Hydra I was either quite massive or has been significantly heated over time."
+" Based on velocity dispersion, Schlaufmanetal. havesuggested that the progenitors of their ECHOS (2011)were dwarf spheroidal galaxies."," Based on velocity dispersion, \citet{schlaufman2011} havesuggested that the progenitors of their ECHOS were dwarf spheroidal galaxies."
+" However, based on the the orbit constraints"," However, based on the the orbit constraints"
+Radiative physics is crucial to the simulation of any astrophysical fluid.,Radiative physics is crucial to the simulation of any astrophysical fluid.
+ Radiative transfer in SPH has a long history. with efforts ranging from simple parametrisations for the radiative cooling time (e.g. ?2)). optical-depth dependent radiative cooling (e.g. 2)) through to flux-limited diffusion models (e.g. 222?222)).," Radiative transfer in SPH has a long history, with efforts ranging from simple parametrisations for the radiative cooling time (e.g. \citealt{Ken_1}) ), optical-depth dependent radiative cooling (e.g. \citealt{Stam_2007}) ), through to flux-limited diffusion models (e.g. \citealt{WB_1,Bastien_diffusion,Viau_et_al_06,WB_2, Mayer_et_al_07, intro_hybrid}) )."
+ While these algorithms are extremely well-suited to calculating gas temperatures during runtime. they lack the precision and insight offered by radiative transfer techniques which do not average over frequency. such as MCRT.," While these algorithms are extremely well-suited to calculating gas temperatures during runtime, they lack the precision and insight offered by radiative transfer techniques which do not average over frequency, such as MCRT."
+ MCRT techniques can be applied to the SPH density field. and have been on several occasions.," MCRT techniques can be applied to the SPH density field, and have been on several occasions."
+ Early attempts began by binning the particle distribution onto a grid - however. the choice of geometry strongly influences the tinal gridded field. and often adaptive meshes are required to correctly represent the matter distribution (2???)..," Early attempts began by binning the particle distribution onto a grid - however, the choice of geometry strongly influences the final gridded field, and often adaptive meshes are required to correctly represent the matter distribution \citep{Oxley_Woolfson_2003,Kurosawa_et_al_04,Stam_MCRT,Bisbas_et_al_09}."
+ Later efforts have utilised ray-tracing directly in SPH fields (???).. allowing the full power of the SPH formalism to be applied to the calculation of optical depths.," Later efforts have utilised ray-tracing directly in SPH fields \citep{Kessel_Deynet_and_Burkert_00,SPHRAY,TRAPHIC}, allowing the full power of the SPH formalism to be applied to the calculation of optical depths."
+ However. these methods currently only calculate optical depths along the ray for the purpose of photoionisation. etc. and do not account for detailed scattering and polarisation. whieh are important in imaging scale systems such as protoplanetary dises.," However, these methods currently only calculate optical depths along the ray for the purpose of photoionisation, etc, and do not account for detailed scattering and polarisation, which are important in imaging small-scale systems such as protoplanetary discs."
+ This paper introduces an algorithm. for imaging SPH simulations of star and planet formation directly using MCRT. without requiring gridding. and including scattering and polarisation.," This paper introduces an algorithm for imaging SPH simulations of star and planet formation directly using MCRT, without requiring gridding, and including scattering and polarisation."
+ It utilises the same techniques that a traditional gridded MCRT code uses. with the advantage that it can trace rays in the density field with the same adaptive capability as SPH. such that it can model radiative effects with at least the same resolving power.," It utilises the same techniques that a traditional gridded MCRT code uses, with the advantage that it can trace rays in the density field with the same adaptive capability as SPH, such that it can model radiative effects with at least the same resolving power."
+ The paper is organised as follows: section 2. will outline the algorithm. sections ?2?. and 4 will describe some applications of the technique to imaging numerical simulations. and section 2? will summarise the results of this work.," The paper is organised as follows: section \ref{sec:method} will outline the algorithm, sections \ref{sec:applications} and 4 will describe some applications of the technique to imaging numerical simulations, and section \ref{sec:conclusions} will summarise the results of this work."
+ For completeness. a summary of MCRT is given here.," For completeness, a summary of MCRT is given here."
+ are emitted from sources (these can be point sources or diffuse emission from the density field itself)., are emitted from sources (these can be point sources or diffuse emission from the density field itself).
+ These packets (hereafter referred to simply as photons) then travel through the medium. and interact with the medium stochastically. reproducing the statistical scattering properties of the medium.," These packets (hereafter referred to simply as photons) then travel through the medium, and interact with the medium stochastically, reproducing the statistical scattering properties of the medium."
+ The exact formalism is summarised below., The exact formalism is summarised below.
+ Photons can be emitted either from a central source (e.g. a protostar). or as diffuse emission from the density field itself. which is only done if the temperature structure of the diffuse component is known (e.g. as output from a simulation).," Photons can be emitted either from a central source (e.g. a protostar), or as diffuse emission from the density field itself, which is only done if the temperature structure of the diffuse component is known (e.g. as output from a simulation)."
+ methods in MCRT (e.g. 2)) can calculate the temperature structure directly from any luminosity source (e.g. stellar. accretion. external radiation fields).," methods in MCRT (e.g. \citealt{Bjorkman_Wood_MCRE}) ) can calculate the temperature structure directly from any luminosity source (e.g. stellar, accretion, external radiation fields)."
+ However. as the SPH simulations used in this work produce temperatures for each SPH particle self-consistently. the system is assumed to already be in temperature equilibrium.," However, as the SPH simulations used in this work produce temperatures for each SPH particle self-consistently, the system is assumed to already be in temperature equilibrium."
+" Tt is also assumed that the dust is thermally coupled to the gas. and that 775,Tistiatien."," It is also assumed that the dust is thermally coupled to the gas, and that $T_{dust} = T_{radiation}$."
+ Tus is suitable for most purposes. except where significant stellar irradiation dominates the radiation field (not modelled in the simulations described in this paper).," This is suitable for most purposes, except where significant stellar irradiation dominates the radiation field (not modelled in the simulations described in this paper)."
+ However. as has been already saic. radiative equilibrium techniques can be used to sidestep this issue: an implementation of radiative equilibrium in SPH fields is discussed in a later section.," However, as has been already said, radiative equilibrium techniques can be used to sidestep this issue: an implementation of radiative equilibrium in SPH fields is discussed in a later section."
+" It is reasonable to (initially) assume all objects emit according to a blackbody spectrum. which then gives: for source and diffuse emission respectively (given the frequency Dye of interest). where c,(r) is the emissivity interpolated from the SPH particle field. i.e. where the sum over j indicates a sum over nearest neighbours. αν is the absorptive opacity. 7; 1s the temperature of each SPH particle. mj and p; are the masses and densities respectively. and W is the smoothing kernel (see section 2.2.1. for more detail on the SPH formalism)."," It is reasonable to (initially) assume all objects emit according to a blackbody spectrum, which then gives: for source and diffuse emission respectively (given the frequency range, of interest), where $\epsilon_{\nu} (\mathbf{r})$ is the emissivity interpolated from the SPH particle field, i.e. where the sum over $j$ indicates a sum over nearest neighbours, $\kappa_{\nu}$ is the absorptive opacity, $T_j$ is the temperature of each SPH particle, $m_j$ and $\rho_j$ are the masses and densities respectively, and $\mathbf{W}$ is the smoothing kernel (see section \ref{sec:tau_SPH} for more detail on the SPH formalism)."
+ There is an extra subtlety regarding the frequency distribution of photons emitted from the gas., There is an extra subtlety regarding the frequency distribution of photons emitted from the gas.
+ The frequency distribution will depend on the local emissivity. which as shown above is an interpolative sum.," The frequency distribution will depend on the local emissivity, which as shown above is an interpolative sum."
+ While in theory the full form of equation (20) should be used to calculate gas luminosity (and the frequency distribution of the photons it emits). this paper approximates the sum using the contribution from one particle only. i.e. This saves computational expense. and is sufficiently accurate for the examples in this paper. bearing in mind a) that the approximation breaks down only in the inner regions of the disc. which are already under-resolved for other reasons (see section 4.23) and b) the effective resolution of the images is too low for this effect to be significant.," While in theory the full form of equation \ref{eq:L_diffuse}) ) should be used to calculate gas luminosity (and the frequency distribution of the photons it emits), this paper approximates the sum using the contribution from one particle only, i.e. This saves computational expense, and is sufficiently accurate for the examples in this paper, bearing in mind a) that the approximation breaks down only in the inner regions of the disc, which are already under-resolved for other reasons (see section \ref{sec:res}) ) and b) the effective resolution of the images is too low for this effect to be significant."
+ The luminosity of each object (whether pointmass or SPH particle) detines how many photon packets are emitted from that object using In order to correctly reproduce the scattering of photons in the medium. the cumulative distribution function (CDF) of interaction must be reproduced.," The luminosity of each object (whether pointmass or SPH particle) defines how many photon packets are emitted from that object using In order to correctly reproduce the scattering of photons in the medium, the cumulative distribution function (CDF) of interaction must be reproduced."
+ In essence. this implies sampling an optical depth using where Visarandomnumberbel wcenQand.," In essence, this implies sampling an optical depth using where is a random number between 0 and 1."
+ LhephotonwilltracetadistanceLL along a constant vector path (hereafter rav)). such that," The photon will travel a distance along a constant vector path (hereafter ), such that"
+"reionization and Umaz= 50kms!, and patchy reionization and Umar=30kms! respectively.","reionization and $v_{max}=50 \mbox{ km
+s}^{-1}$, and patchy reionization and $v_{max}=30 \mbox{ km s}^{-1}$ respectively."
+" The three lower panels show the same, but with luminosities determined by abundance matching (eqn. 1))."," The three lower panels show the same, but with luminosities determined by abundance matching (eqn. \ref{eq:ab_match}) )."
+" While the models with post-reionization star formation suppressed at Όπιας=50kms! give a good match to the observed luminosity function, when we decrease the suppression threshold to 30kms! none of the models fit the bias-corrected luminosity function well, despite the overall number of satellites only being slightly increased."," While the models with post-reionization star formation suppressed at $v_{max}=50 \mbox{ km s}^{-1}$ give a good match to the observed luminosity function, when we decrease the suppression threshold to $30
+\mbox{ km s}^{-1}$ none of the models fit the bias-corrected luminosity function well, despite the overall number of satellites only being slightly increased."
+" This holds true both for the abundance matching-derived luminosities, and the best-fit SFR/SPS model."," This holds true both for the abundance matching-derived luminosities, and the best-fit SFR/SPS model."
+" In particular, even if the star formation efficiency is decreased to give the correct number of satellites overall, the shape still does not match the observed luminosity function."," In particular, even if the star formation efficiency is decreased to give the correct number of satellites overall, the shape still does not match the observed luminosity function."
+" It is possible, however, that this lower suppression model could still be a good match, if we explicitly consider other processes like supernova feedback (Lietal.2010).."," It is possible, however, that this lower suppression model could still be a good match, if we explicitly consider other processes like supernova feedback \citep{Li2010}."
+" Compared to patchy reionization, the instantaneous models generally yield more satellites."," Compared to patchy reionization, the instantaneous models generally yield more satellites."
+" While only the D halo is shown here as an illustration, as discussed in Section 4.1 this holds true for all of the Aquarius halos, except for the F halo."," While only the D halo is shown here as an illustration, as discussed in Section \ref{sec:reion_res} this holds true for all of the Aquarius halos, except for the F halo."
+" It is interesting to note, however, that the differences between different reionization models, luminosity prescriptions, etc.,"," It is interesting to note, however, that the differences between different reionization models, luminosity prescriptions, etc.,"
+ are in general dominated by the halo-halo scatter., are in general dominated by the halo-halo scatter.
+" That is, the differences in the satellite populations of the six halos for any given model are as large as the differences between the range of models we are exploring."," That is, the differences in the satellite populations of the six halos for any given model are as large as the differences between the range of models we are exploring."
+" As an example, Figure 10 shows the satellite luminosity functions for all six halos overplotted, here for reionization model 1 and the SFR luminosity prescription."," As an example, Figure \ref{fig:all_halos} shows the satellite luminosity functions for all six halos overplotted, here for reionization model 1 and the SFR luminosity prescription."
+" While the B halo, which is the clear outlier, is also the lowest mass halo in this simulation suite, the uncertainty of where the Milky Way itself fits in this distribution makes direct comparisons to constrain models difficult."," While the B halo, which is the clear outlier, is also the lowest mass halo in this simulation suite, the uncertainty of where the Milky Way itself fits in this distribution makes direct comparisons to constrain models difficult."
+" In our models, we find that the population of faint satellites is highly sensitive to the reionization model."," In our models, we find that the population of faint satellites is highly sensitive to the reionization model."
+ This agrees with the results of," This agrees with the results of \citet{Busha2010, Font2011},"
+starburst.,starburst.
+ Assuming that most of the energv in photons emitted by voung stars eventually results in heating the cust (svhieh then re-radiates it as black-body thermal continuum in photons). (hen star formation rates for stars more massive Chan 5M. can be estimated from the far-IR. luminosity (Cram et al 1993).," Assuming that most of the energy in photons emitted by young stars eventually results in heating the dust (which then re-radiates it as black-body thermal continuum in photons), then star formation rates for stars more massive than $_\odot$ can be estimated from the far-IR luminosity (Cram et al 1998)."
+ Such stars dominate the total stellar luminosity. and end (heir lives as supernovae.," Such stars dominate the total stellar luminosity, and end their lives as supernovae."
+ Supernova rates in starbursts can be cirectly related io formation rates of massive stars. ancl by assuming an initial mass [function (IME). the luminosities of the starbursts.," Supernova rates in starbursts can be directly related to formation rates of massive stars, and by assuming an initial mass function (IMF), the luminosities of the starbursts."
+ SGudies of radio supernovae (RSNe) in starbursts are unlike studies in our own Galaxy., Studies of radio supernovae (RSNe) in starbursts are unlike studies in our own Galaxy.
+ The sample of RSNe in a starburst is well-defined and largely [ree from selection effects. ancl all the sources can be considered to be at the same distance.," The sample of RSNe in a starburst is well-defined and largely free from selection effects, and all the sources can be considered to be at the same distance."
+ Therefore. nearby. starbursts can be considered as laboratories in which to investigate. both statistically and individually. the properties of a well-defined sample of RSNe and supernova remnants (SNR).," Therefore, nearby starbursts can be considered as laboratories in which to investigate, both statistically and individually, the properties of a well-defined sample of RSNe and supernova remnants (SNR)."
+ Arp 220. at a distance of T7Mpe and with an infrared Iuminosityv of Lye1.3x10151. (Soiler et al..," Arp 220, at a distance of 77Mpc and with an infrared luminosity of $_{IR} \sim 1.3 \times 10^{12}$ $_\odot$ (Soifer et al.,"
+ 1937) is one of the closest and best studied Ultra Luminous Inlralted Galaxies. or ULIRGs.," 1987) is one of the closest and best studied Ultra Luminous InfraRed Galaxies, or ULIRGs."
+ It is. like most ULIBRGs a merger of two gas-rich galaxies: Arp 220 has a pair of radio nuclei separated in roughly (the E-W direction by approximately 370 pe (Norris. 1983. Rovilos et al. 2003. Rocrignez-Rico et al..," It is, like most ULIRGs a merger of two gas-rich galaxies; Arp 220 has a pair of radio nuclei separated in roughly the E-W direction by approximately 370 pc (Norris, 1988, Rovilos et al, 2003, Rodriguez-Rico et al.,"
+ 2005)., 2005).
+ The two nuclei are also clearly seen by high-resolution Keck imaging in the mid-infrared (Soifer et al..," The two nuclei are also clearly seen by high-resolution Keck imaging in the mid-infrared (Soifer et al.,"
+ 1999. 2002).," 1999, 2002)."
+ The western nucleus is brighter than (lie eastern one. though Che flux ratio is Irequeney. dependent.," The western nucleus is brighter than the eastern one, though the flux ratio is frequency dependent."
+ Smith et al (1993a) made 1δο continuum observations of Arp 220. in which thev found a number of unresolved sources in (he western nucleus which thev interpreted. as huninous radio supernovae.," Smith et al (1998a) made 18cm continuum observations of Arp 220, in which they found a number of unresolved sources in the western nucleus which they interpreted as luminous radio supernovae."
+ Thev suggested (hat this interpretation is consistent with a simple starburst model for the IR huninosity of Arp 220 that has a star formation rate of 50-100. AL. /vr and a luninous supernova rate of 1.75-3.5/vr., They suggested that this interpretation is consistent with a simple starburst model for the IR luminosity of Arp 220 that has a star formation rate of 50-100 $_\odot$ /yr and a luminous supernova rate of 1.75-3.5/yr.
+ The supernovae observed are the most clearly detected. RSNe in a distant galaxy., The supernovae observed are the most clearly detected RSNe in a distant galaxy.
+ Smith et al suggested. based on the extreme [παν densities. (hat the objects mieht be Type IIn supernovae.," Smith et al suggested, based on the extreme flux densities, that the objects might be Type IIn supernovae."
+ This twpe οἱ supernova is oplically identilied by a narrow component in its spectrum and is thought to be interacting with a dense circumstellar medium. a result of the high mass loss rates of the progenitor stars.," This type of supernova is optically identified by a narrow component in its spectrum and is thought to be interacting with a dense circumstellar medium, a result of the high mass loss rates of the progenitor stars."
+ Rovilos et al (2005) obtained multi-epoeh high resolution data of Arp 220. and concluded that the light curves had been surprisingly stable over a period of more (han five vears.," Rovilos et al (2005) obtained multi-epoch high resolution data of Arp 220, and concluded that the light curves had been surprisingly stable over a period of more than five years."
+ The long decay time of the 13cm light curves of these sources does not agree with the simple model of Type Hn supernovae (Smith et al., The long decay time of the 18cm light curves of these sources does not agree with the simple model of Type IIn supernovae (Smith et al.
+ 1998a)., 1998a).
+ The observations discussed in this paper are part of a subsequent high-sensitivity monitoring programme of Arp220., The observations discussed in this paper are part of a subsequent high-sensitivity monitoring programme of Arp220.
+"Active galactic nuclei (AGN) have been studied extensively in the wavelength range from em-radio to y-rays in the last decades (see. e.g.. Kembhavi Narlikar 1900, or Krolik 1999., and references therein for a review).","Active galactic nuclei (AGN) have been studied extensively in the wavelength range from cm-radio to $\gamma$ -rays in the last decades (see, e.g., Kembhavi Narlikar \cite{kembhavi1999}, or Krolik \cite{krolik1999}, and references therein for a review)."
+" There is overwhelming observational evidence that their emission originates from accretion onto supermassive black holes (SMBH) with masses M,=109-?M.. (e.g.. Ferrarese Ford 2005.. and references therein)."," There is overwhelming observational evidence that their emission originates from accretion onto supermassive black holes (SMBH) with masses $M_{\bullet}\simeq10^{6...9}M_{\odot}$ (e.g., Ferrarese Ford \cite{ferrarese2005}, and references therein)."
+ The viewing-angle unification standard model of AGN (Lawrence 1987:: see also Urry Padovani 1995)) has long been established. although recent studies emphasize the impact of galaxy mergers on AGN fueling and luminosity (e.g. Hopkins et al. 2006)).," The viewing-angle unification standard model of AGN (Lawrence \cite{lawrence1987}; see also Urry Padovani \cite{urry1995}) ) has long been established, although recent studies emphasize the impact of galaxy mergers on AGN fueling and luminosity (e.g. Hopkins et al. \cite{hopkins2006}) )."
+ However. the details of AGN activity are still not well understood.," However, the details of AGN activity are still not well understood."
+ One example is the nuclear emission: proposals for its origin range from shocks in continuous (e.g.. Marscher Gear 1985)) or discontinuous (e.g.. Spada et al. 2001))," One example is the nuclear emission; proposals for its origin range from shocks in continuous (e.g., Marscher Gear \cite{marscher1985}) ) or discontinuous (e.g., Spada et al. \cite{spada2001}) )"
+" Jets to orbiting plasma ""hotspots"" (e.g.. Abramowiez et al. 1991))"," jets to orbiting plasma “hotspots” (e.g., Abramowicz et al. \cite{abramowicz1991}) )"
+ or plasma density waves (e.g.. Kato 2000)) in accretion disks.," or plasma density waves (e.g., Kato \cite{kato2000}) ) in accretion disks."
+ One possibility to constrain the AGN physics is offered by the analysis of characteristic timescales and the noise properties of the emission., One possibility to constrain the AGN physics is offered by the analysis of characteristic timescales and the noise properties of the emission.
+ Several studies aimed at identifying characteristic timescales. including possible quasi-periodic oscillations. from tens of minutes (using X-ray observations: e.g. Benlloch et al. 2001..," Several studies aimed at identifying characteristic timescales, including possible quasi-periodic oscillations, from tens of minutes (using X-ray observations; e.g. Benlloch et al. \cite{benlloch2001},"
+ and using mm-radio observations: e.g. Schóddel et al. 2007)), and using mm-radio observations; e.g. Schöddel et al. \cite{schoedel2007}) )
+ to tens of years (based on em- and mm-radio monitoring data: Hovatta et al. 2007. 2008)).," to tens of years (based on cm- and mm-radio monitoring data; Hovatta et al. \cite{hovatta2007,hovatta2008}) )."
+ Other works focused on the power spectral statistics of the emission which was found to globally follow red-noise type powerlaws (e.g. Lawrence Papadakis 1993:;; Uttley et al. 2002))., Other works focused on the power spectral statistics of the emission which was found to globally follow red-noise type powerlaws (e.g. Lawrence Papadakis \cite{lawrence1993}; Uttley et al. \cite{uttley2002}) ).
+ The studies by Hovatta et al. (2007. 2008))," The studies by Hovatta et al. \cite{hovatta2007,hovatta2008}) )"
+ used radio data in the frequency range 22-87 GHz spanning timelines of x25 years., used radio data in the frequency range 22–87 GHz spanning timelines of $\approx$ 25 years.
+ A key result of their study was that AGN show large-amplitude activity on timescales z6 years., A key result of their study was that AGN show large-amplitude activity on timescales $\gtrsim$ 6 years.
+ Therefore. analysing the long-term behaviour of AGN requires monitoring on timescales of 510 years.," Therefore, analysing the long-term behaviour of AGN requires monitoring on timescales of $\gtrsim$ 10 years."
+ Building up on this conclusion. we have analyzed archival flux monitoring data from the IRAM Plateau de Bure Interferometer (PdBI) of six AGN 09234392. 3C Ill. 3C 273. 3C 345. 3C 454.3. and 3C 84 - in the (observatory frame) frequency range 80-267 GHz.," Building up on this conclusion, we have analyzed archival flux monitoring data from the IRAM Plateau de Bure Interferometer (PdBI) of six AGN – 0923+392, 3C 111, 3C 273, 3C 345, 3C 454.3, and 3C 84 – in the (observatory frame) frequency range 80--267 GHz."
+ Given the range of redshifts ο=0.02—0.86 of our sample. this corresponds to a range of rest-frame frequencies of 82-497 GHz.," Given the range of redshifts $z\approx0.02-0.86$ of our sample, this corresponds to a range of rest-frame frequencies of $\approx$ 82–497 GHz."
+ Our dataset comprises several thousand measurements for each source. covering. timelines of x14 years.," Our dataset comprises several thousand measurements for each source, covering timelines of $\approx$ 14 years."
+" We note that our study complements earlier AGN mmyradio surveys carried out with the IRAM Pico Veleta observatory (Steppe et al. 1988.1992,1993;"," We note that our study complements earlier AGN mm/radio surveys carried out with the IRAM Pico Veleta observatory (Steppe et al. \cite{steppe1988,steppe1992,steppe1993};"
+ Reuter et al. 1997))., Reuter et al. \cite{reuter1997}) ).
+ The PdBI is able to observe in three atmospheric. windows located around wavelengths of 1.3 mm. 2 mm. and 3 mm.," The PdBI is able to observe in three atmospheric windows located around wavelengths of 1.3 mm, 2 mm, and 3 mm."
+ Each of these bands covers a continuous range of frequencies that are available for observations: these ranges are 201-267 GHz for the 1.3 mm band. 129-174 GHz for the 2 mm band. and 80-116 GHz for the 3 mm band (Winters Neri 2010)).," Each of these bands covers a continuous range of frequencies that are available for observations; these ranges are 201–267 GHz for the 1.3 mm band, 129–174 GHz for the 2 mm band, and 80–116 GHz for the 3 mm band (Winters Neri \cite{winters2010}) )."
+ Until the end of 2006. observations were made using dual-band. single-polarization receivers.," Until the end of 2006, observations were made using dual-band, single-polarization receivers."
+ One frequency in the 1.3- band and one frequency in the 3-mm band could be selected for simultaneous observations; the 2-mm band was not reachable., One frequency in the 1.3-mm band and one frequency in the 3-mm band could be selected for simultaneous observations; the 2-mm band was not reachable.
+ At the beginning of 2007. new single-band. dual-polarization receivers were installed.," At the beginning of 2007, new single-band, dual-polarization receivers were installed."
+ Since then it is possible to observe a single frequency located in a given band simultaneously in both linear polarizations., Since then it is possible to observe a single frequency located in a given band simultaneously in both linear polarizations.
+ At the beginning of 2008. the 2-mm band became available for observations.," At the beginning of 2008, the 2-mm band became available for observations."
+ Flux monitoring data for AGN are available from October 1996 on as a by-product of regular observatory operations., Flux monitoring data for AGN are available from October 1996 on as a by-product of regular observatory operations.
+" The PdBI uses the most luminous (S,z3 Jy in the", The PdBI uses the most luminous $S_{\nu}\gtrsim3$ Jy in the
+504. we compare their column densities predicted by C84 and C07.,", we compare their column densities predicted by C84 and C07."
+ In the previous photoionisation calculations in this paper it was sufficient to use one zone models for optically thin gas., In the previous photoionisation calculations in this paper it was sufficient to use one zone models for optically thin gas.
+ However. to calculate and compare the column densities of the ions over the range 1.0«log£<4.5. we chose to specify the gas to have total hydrogen column density Vy=1077em.7. which is typical for warm absorbers.," However, to calculate and compare the column densities of the ions over the range $1.0<\log\xi<4.5$, we chose to specify the gas to have total hydrogen column density $\nh=\onlyten{22}\cmsqi$, which is typical for warm absorbers."
+ The column densities are plotted in with solid lines for CO7 and dashed lines for C84., The column densities are plotted in with solid lines for C07 and dashed lines for C84.
+ It is seen that the column densities of the major coolants changed significantly., It is seen that the column densities of the major coolants changed significantly.
+ The cooling agents are among the ions for which dielectronic recombination rate coefficients (hereafter DRRC) have been updated to the references below. as will be discussed in detail in4.," The cooling agents are among the ions for which dielectronic recombination rate coefficients (hereafter DRRC) have been updated to the references below, as will be discussed in detail in."
+. Thus the enhanced cooling in CO7 due to the change in DRRC is the cause of the shift in the stability curves., Thus the enhanced cooling in C07 due to the change in DRRC is the cause of the shift in the stability curves.
+" The evolution of the thermal phases from the nebular temperatures of 107Ix to the coronal temperatures of LO""Ix depends sensitively on the detailed atomic physics of the various elements which contribute to photoelectric heating as well as to cooling due to recombination and collisionally excited lines.", The evolution of the thermal phases from the nebular temperatures of $10^4\kel$ to the coronal temperatures of $10^6\kel$ depends sensitively on the detailed atomic physics of the various elements which contribute to photoelectric heating as well as to cooling due to recombination and collisionally excited lines.
+ In2.. we have compared the total recombination rates (dielectronic + radiative) for the significant cooling agents as predicted by CO7 and C84.," In, we have compared the total recombination rates (dielectronic + radiative) for the significant cooling agents as predicted by C07 and C84."
+ The values for log Zand log£ at which the comparisons have been made are given in columns 2 and 3 respectively., The values for $\log T$ and $\log \xi$ at which the comparisons have been made are given in columns 2 and 3 respectively.
+ For all the ions, For all the ions
+over-luminous relative to the mean relation.,over-luminous relative to the mean relation.
+ However. this ellect is very small ancl such populations. --[ present. must be weaker than in most field ancl cluster ellipticals. given the ugh degree of homogeneity in BCC Ομ. values.," However, this effect is very small and such populations, if present, must be weaker than in most field and cluster ellipticals, given the high degree of homogeneity in BCG $_{EW}$ values."
+ ]t may be possible to use these correlations to celine uigher precision distance indicators for BCCs. which could »e used for galaxies with redshifts up to 18.000. as is indicated by the reduced. scatter in. ραπ absolute magnitude. and the reduced. scatter about the structure »wameter relation after correction for the correlation with gi. discussed in section 3.," It may be possible to use these correlations to define higher precision distance indicators for BCGs, which could be used for galaxies with redshifts up to 18,000 $^{-1}$, as is indicated by the reduced scatter in R-band absolute magnitude, and the reduced scatter about the structure parameter relation after correction for the correlation with $_{EW}$, discussed in section 3."
+ Recent ROSAT observations reveal that a significant number of the Lauer Postman (1994). galaxies do not lie at the X-ray centroids of their clusters (Paul Lvnam. private communication). and there may. be better candidates: for the dominant central cluster galaxy.," Recent ROSAT observations reveal that a significant number of the Lauer Postman \shortcite{la:94} galaxies do not lie at the X-ray centroids of their clusters (Paul Lynam, private communication), and there may be better candidates for the dominant central cluster galaxy."
+ Thus our conclusions may refer more generally to bright galaxies towards cluster centres than to individual DCCs inhabiting the very centre of the cluster potential., Thus our conclusions may refer more generally to bright galaxies towards cluster centres than to individual BCGs inhabiting the very centre of the cluster potential.
+ We acknowledge the anonymous referee. for many useful suggestions., We acknowledge the anonymous referee for many useful suggestions.
+ The United. Winecom Infrared. Telescope is operated. by the Joint Astronomy Centre on behalf of the Ulx Particle Physics and Astronomy Research Council., The United Kingdom Infrared Telescope is operated by the Joint Astronomy Centre on behalf of the UK Particle Physics and Astronomy Research Council.
+"Here p and v stand for density and velocity of the gas, Pc and Pg pressures of the CRs and the thermal gas correspondingly.","Here $\rho$ and $v$ stand for density and velocity of the gas, $P_C$ and $P_G$ pressures of the CRs and the thermal gas correspondingly."
+ This system of equations describes mass and momentum conservation (first two equations) and entropy conservation for gas and cosmic rays (last two equations)., This system of equations describes mass and momentum conservation (first two equations) and entropy conservation for gas and cosmic rays (last two equations).
+" The boundary conditions come from applying generalized (i.e. cosmic rays included) jump conditions at the shock, which come from mass, momentum and energy conservation for particles that cross the shock, see ? for details."," The boundary conditions come from applying generalized (i.e. cosmic rays included) jump conditions at the shock, which come from mass, momentum and energy conservation for particles that cross the shock, see \cite{Chevalier1983} for details."
+" To preserve self-similarity, the shock compression ratio, s, and the proportion of the total pressure going to cosmic rays w (CR acceleration efficiency) must be taken as constants giving: In fact s and w should be taken from a nonlinear shock model, but at the level of this model it is enough to treat them as adjustable parameters."," To preserve self-similarity, the shock compression ratio, $s$, and the proportion of the total pressure going to cosmic rays $w$ (CR acceleration efficiency) must be taken as constants giving: In fact $s$ and $w$ should be taken from a nonlinear shock model, but at the level of this model it is enough to treat them as adjustable parameters."
+" In standard fashion the system of partial differential equations (I)) can now be reduced to a system of ordinary differential equations by transforming to the similarity variable €οςrt7/5, with the constant of proportionality chosen such such that £=r/rshock(t) (see ? for the constant’s numerical value)."," In standard fashion the system of partial differential equations \ref{eqs-hydro-Chevalier}) ) can now be reduced to a system of ordinary differential equations by transforming to the similarity variable $\xi\propto rt^{-2/5}$, with the constant of proportionality chosen such such that $\xi=r/r_{shock}(t)$ (see \cite{Chevalier1983} for the constant's numerical value)."
+" Introducing scaled structure functions G, U and Z we obtain for the density, velocity, and pressure, respectively, and for the compression ratio at the shock we get wl aul where Ύω—=SQL)_5+3 In the following we will work in the coordinate system defined by the variables € and t, so an element of fluid intersects the shock (£— 1) and then moves with velocity 1) towards the center (£— 0) of the SNR."," Introducing scaled structure functions $G$, $U$ and $Z$ we obtain for the density, velocity, and pressure, respectively, and for the compression ratio at the shock we get $s\equiv\frac{G(\xi=1+)}{G(\xi=1-)}=\frac{\gamma_w+1}{\gamma_w-1}$ , where $\gamma_w=\frac{5+3w}{3(w+1)}$ In the following we will work in the coordinate system defined by the variables $\xi$ and $t$, so an element of fluid intersects the shock $\xi=1$ ) and then moves with velocity $\dot{\xi}=\frac{2}{5}\xi(U(\xi)-1)$ towards the center $\xi=0$ ) of the SNR."
+ The solution of these equations for different w are shown in Fig[I]. where n—p/my is the particle number density.," The solution of these equations for different $w$ are shown in \ref{fig:rho-v-p-chevalier}, where $n=\rho/m_p$ is the particle number density."
+ We now consider Coulomb electron interactions with hot protons in the SNR model described in the previous subsection., We now consider Coulomb electron interactions with hot protons in the SNR model described in the previous subsection.
+" In the coordinate system that moves with the plasma, the equation that describes electron heating, can be found e.g. in"," In the coordinate system that moves with the plasma, the equation that describes electron heating, can be found e.g. in"
+There is also ambiguity in the value of the parameter sc in the definition of ry [eq. (5))].,There is also ambiguity in the value of the parameter $\mu$ in the definition of $r_T$ [eq. \ref{rT}) )].
+ Numerical simulations show that je is usually of order unity. the value we use. but that it can Γje as larelarge as L.71.7 in specilicspecili cases (Frol(Frolovet{al.1994:Dienerοἱ1997).," Numerical simulations show that $\mu$ is usually of order unity, the value we use, but that it can be as large as $1.7$ in specific cases \citep{fro94,die97}."
+.1997).. The time intervalinterva Aly [eq. (9))], The time interval $\Delta t_1$ [eq. \ref{t1}) )]
+ is very sensitive to the value of jc Alyx40.," is very sensitive to the value of $\mu$ : $\Delta t_1
+\propto \mu^3$ ."
+" Because of the tight observational constraint on Al, (Figs.", Because of the tight observational constraint on $\Delta t_1$ (Figs.
+ 2 and 3)). pe cannot be much larger than | since otherwise it would imply too small a black hole mass.," \ref{fig2} and \ref{fig3}) ), $\mu$ cannot be much larger than $1$ since otherwise it would imply too small a black hole mass."
+" Another constraint on Mj; aud. M, is obtained from the requirement that the radius rg: al which the star is (ically disrupted must lie outside the radius of the marginally bound circular orbit.", Another constraint on $M_H$ and $M_\star$ is obtained from the requirement that the radius $r_T$ at which the star is tidally disrupted must lie outside the radius of the marginally bound circular orbit.
+" This requires ry>tr, for a Schwarzschild black hole and ry> for the most favorably oriented orbit around a maximally spinning Kerr hole.", This requires $r_T>4r_g$ for a Schwarzschild black hole and $r_T>r_g$ for the most favorably oriented orbit around a maximally spinning Kerr hole.
+ The uncertainty in the limiting rp is indicated by the dotted lines in Figure 2.., The uncertainty in the limiting $r_T$ is indicated by the dotted lines in Figure \ref{fig2}.
+" For concreteness. we have assumect that the limiting radius is intermediate between the (wo extreme cases mentioned above: we have used the limit ry>2r,."," For concreteness, we have assumed that the limiting radius is intermediate between the two extreme cases mentioned above; we have used the limit $r_T>2r_g$."
+ Yet another source of uncertainty derives from the [act that we have used Newtonian dynamics in our analysis. whereas it is clear (hat relativistic effects should be important (Luminet&Marck1955:Lagunaetal.1993:οἱal. 1997).," Yet another source of uncertainty derives from the fact that we have used Newtonian dynamics in our analysis, whereas it is clear that relativistic effects should be important \citep{lum85,lag93,fro94,die97}."
+. lgnoring these caveats. we obtain a üght correlation between (he mass of the disrupted star in NGC! 5905 and the mass of the central black hole.," Ignoring these caveats, we obtain a tight correlation between the mass of the disrupted star in NGC 5905 and the mass of the central black hole."
+ As Figure 2. shows. the data are consistent. with two distinct solution branches.," As Figure \ref{fig2} shows, the data are consistent with two distinct solution branches."
+" One solution branch corresponds to a relatively massive black hole. Mj; close to 105M... lidally disruptiing a main sequence star wilh AZ,~.M.."," One solution branch corresponds to a relatively massive black hole, $M_H$ close to $10^8M_\odot$, tidally disrupting a main sequence star with $M_\star\sim M_\odot$."
+ The other branch consists of a lower mass black hole. Mj;<LO'OAL.. disrupting a brown dwarf or a giant planet with AL0.02.4...," The other branch consists of a lower mass black hole, $M_H < 10^{7.5}M_\odot$, disrupting a brown dwarf or a giant planet with $M_\star \la 0.02M_\odot$."
+" Assuming that NGC 5905 satisfies the well-known correlations between black hole mass and host galaxy properties (Saluceietal.2000;GebhardtFerrarese 2001).. the black hole mass is estimated to be Mj,10*—LOSAL. (IXomossa2001)."," Assuming that NGC 5905 satisfies the well-known correlations between black hole mass and host galaxy properties \citep{sal00,geb00,mer01}, the black hole mass is estimated to be $M_H\sim 10^7-10^8M_\odot$ \citep{kom01}."
+. This mass rauge permits both solution branches., This mass range permits both solution branches.
+ Figures 4 and 5 show sample theoretical light curves corresponding to the two solution branches., Figures \ref{fig4} and \ref{fig5} show sample theoretical light curves corresponding to the two solution branches.
+ In each case. there is one [ree parameter. namely the Iraction / of the mass of the disrupted star that returns as fallback debris to produce the observed radiation.," In each case, there is one free parameter, namely the fraction $f$ of the mass of the disrupted star that returns as fallback debris to produce the observed radiation."
+ We find that f is unusually small if the clisrupted star is a main sequence star: lor the particular example shown in Figure 4.. we require /«0.0003. which is an extremely small value compared to thetheoretically. expected value of f[~0.1—0.5.," We find that $f$ is unusually small if the disrupted star is a main sequence star; for the particular example shown in Figure \ref{fig4}, we require $f<0.0003$, which is an extremely small value compared to thetheoretically expected value of $f\sim 0.1-0.5$."
+ Brown cwarf models give somewhat larger estimates lor f (Fig. 5)).," Brown dwarf models give somewhat larger estimates for $f$ (Fig. \ref{fig5}) ),"
+ though still uncomfortably small. unless the black hole has an umisuallv lowmass (<105?M. ).," though still uncomfortably small, unless the black hole has an unusually lowmass $<10^{6.5}M_\odot$ )."
+ Though the estimated / can be increased by adopting a larger ambiguity factor jin the definition of rp. we do not consider this likely sincea larger je vill lead to an unacceptably small value of Mj. ," Though the estimated $f$ can be increased by adopting a larger ambiguity factor $\mu$ in the definition of $r_T$ , we do not consider this likely sincea larger $\mu$ will lead to an unacceptably small value of $M_H.$ "
+ , 
+latter corresponds to a magnetized cireumburst medium.,latter corresponds to a magnetized circumburst medium.
+ It is this possibility that will be addressed and tested in the present work., It is this possibility that will be addressed and tested in the present work.
+ Section 3 presents an analytical discussion of the model and is followed by numerical calculations of the light curve in Sec. ??.., Section \ref{sec:model} presents an analytical discussion of the model and is followed by numerical calculations of the light curve in Sec. \ref{sec:lcurve}. .
+ Section 2? discusses the results in the light of modern X-ray afterglows., Section \ref{sec:disc} discusses the results in the light of modern X-ray afterglows.
+ We consider the following fiducial values for the parameters characterizing the afterglow., We consider the following fiducial values for the parameters characterizing the afterglow.
+" The ejecta is composed of a homogeneous shell of width A=c with 7= 1071s in the stationary frame. with (isotropic equivalent) bulk Kinetic energy k=107Es erg and Lorentz factor +;=3005,45..."," The ejecta is composed of a homogeneous shell of width $\Delta =c T$ with $T=10T_{1}\,$ s in the stationary frame, with (isotropic equivalent) bulk kinetic energy $E=10^{54}\,E_{54}\,$ erg and Lorentz factor $\gamma_{\rm
+ ej}=300\gamma_{\rm ej,2.5}$."
+" This outflow impinges on a stellar wind with density profile po= Sl with A=510111, g/em and (toroidal) magnetic field profile B.=10°D,G(r107?!: the variables ei, and D, encode our uncertainty on the densitycm) and the magnetic field."," This outflow impinges on a stellar wind with density profile $\rho_{\rm w} = A r^{-2}$ , with $A = 5\times
+10^{11} A_{*}\,$ g/cm and (toroidal) magnetic field profile $B_{\rm
+ w}=10^3B_* \,{\rm G}\, \left(r/10^{12}\,{\rm cm}\right)^{-1}$; the variables $A_*$ and $B_*$ encode our uncertainty on the density and the magnetic field."
+" The value D,=1 corresponds to the formation of a magnetar with surface field ~104° G after the collapse of a 1072.. progenitor star."," The value $B_*=1$ corresponds to the formation of a magnetar with surface field $\sim 10^{15}\,$ G after the collapse of a $10R_\odot$ progenitor star."
+ The magnetic field of Wolf-Rayet stars. which are considered as potential progenitor stars for gamma-ray bursts. are not known. but surface values as large as 10. 10°G have been considered (Ignace et al.," The magnetic field of Wolf-Rayet stars, which are considered as potential progenitor stars for gamma-ray bursts, are not known, but surface values as large as $1-10\times
+10^3\,$ G have been considered (Ignace et al."
+ 1998)., 1998).
+ For these parameters. the magnetization m.=D.(πριc)2L8ϱ1D7 1: it does not depend on r.," For these parameters, the magnetization $\sigma_{\rm w}\equiv B_{\rm
+ w}^2/\left(4\pi \rho_{\rm w}c^2\right)\simeq 1.8\times 10^{-4}
+B_*^2A_*^{-1}$ ; it does not depend on $r$."
+" Note that the central assumption hereA, is that of a relatively high magnetization of the external medium: the density profile does not play a erucial role and similar effects can be observed in a constant density medium of sufficient magnetization. as discussed briefly in Sec. ??.."," Note that the central assumption here is that of a relatively high magnetization of the external medium; the density profile does not play a crucial role and similar effects can be observed in a constant density medium of sufficient magnetization, as discussed briefly in Sec. \ref{sec:disc}."
+" The proper density of the ejecta =La(4i7sAmer ): therefore the density contrast between the ejecta and the external medium is not very large. Ggfm)mSIT,..-ee implying that the reverse shock propagates at relativistic speeds in the ejecta (Sari Piran 1995)."," The proper density of the ejecta $n_{\rm ej} = E_{\rm ej}/\left(4\pi
+ r^2 \gamma_{\rm ej}^2 \Delta m_pc^2\right)$ ; therefore the density contrast between the ejecta and the external medium is not very large, $\left(n_{\rm ej}/n_{\rm w}\right)^{1/2} \simeq 81
+T_{1}^{-1/2}\gamma_{\rm ej,2.5}^{-1}E_{54}^{1/2}A_*^{-1/2}$, implying that the reverse shock propagates at relativistic speeds in the ejecta (Sari Piran 1995)."
+ The shocked material — which we denote as the blast — thus moves with initial Lorentz factor It remains constant as long as the reverse shock is crossing the shell (Sari Piran 1995. Beloborodov Uhm 2006).," The shocked material – which we denote as the blast – thus moves with initial Lorentz factor It remains constant as long as the reverse shock is crossing the shell (Sari Piran 1995, Beloborodov Uhm 2006)."
+ Approximating the velocity of the reverse shock as c in the ejecta frame. the reverse shock has crossed the outflow at radius x.— dn the stationary frame). corresponding to observer time fo=~seell|1. with z the GRB redshift.," Approximating the velocity of the reverse shock as $c$ in the ejecta frame, the reverse shock has crossed the outflow at radius $r_\times =\gamma_{\rm b}^2 c
+T$ (in the stationary frame), corresponding to observer time $t_{\times}\simeq 5\,{\rm sec}\,(1+z)T_{1}$, with $z$ the GRB redshift."
+ Beyond r.. the blast Lorentz factor decreases according to 5n&sn.(rir)L7 dn the adiabatic regime. i. corresponding to Eq. €1))," Beyond $r_\times$, the blast Lorentz factor decreases according to $\gamma_{\rm b}\simeq \gamma_{\rm
+ b,\times}\left(r/r_\times\right)^{-1/2}$ in the adiabatic regime, $\gamma_{\rm b,\times}$ corresponding to Eq. \ref{eq:gammab}) )."
+" The relationship between observer time. radius and blast Lorentz factor then becomes (/,,..//.)GirYemsuf) 7"," The relationship between observer time, radius and blast Lorentz factor then becomes $\left(t_{\rm obs}/t_\times\right) \simeq
+\left(r/r_\times\right)^2 \simeq \left(\gamma_{\rm b}/\gamma_{\rm
+ b,\times}\right)^{-4}$ ."
+ We now come to the modelling of the electron population in the blast., We now come to the modelling of the electron population in the blast.
+" Following Lemoine Pelletier (2010. 2011). we define the parameter dina—&,MEEA which characterizes whether instabilities may develop or not in the shock precursor. hence whether Fermi cycles can develop or not."," Following Lemoine Pelletier (2010, 2011), we define the parameter $Y_{\rm inst} \equiv \xi_{\rm b}^{-1}\sigma_{\rm
+ w}\gamma_{\rm sh}^2$, which characterizes whether instabilities may develop or not in the shock precursor, hence whether Fermi cycles can develop or not."
+" The parameter £1, denotes the fraction of incoming matter energy through the shock τρις that is carried by the accelerated and returning particles (i.e. the beam)."," The parameter $\xi_{\rm b}$ denotes the fraction of incoming matter energy through the shock $4\gamma_{\rm b}^2\rho_{\rm
+ w}c^2$ that is carried by the accelerated and returning particles (i.e. the beam)."
+ By returning. it is meant those incoming protons that are reflected on the shock front. which constitute an essential ingredient of the shock formation.," By returning, it is meant those incoming protons that are reflected on the shock front, which constitute an essential ingredient of the shock formation."
+ These reflected protons exist even in the absence of Fermi powerlaws., These reflected protons exist even in the absence of Fermi powerlaws.
+ Through mixing with he unshocked plasma. these returning particles (along with the accelerated. particles) induce two-stream or filamentation micro-instabilities in the shock precursor. on scales close to the electron o Ion skin depth czσωμα.," Through mixing with the unshocked plasma, these returning particles (along with the accelerated particles) induce two-stream or filamentation micro-instabilities in the shock precursor, on scales close to the electron to ion skin depth $c/\omega_{\rm pe}\rightarrow c/\omega_{\rm pi}$."
+ The filamentation instability has ime to grow only if dius1. while other two stream instabilities may grow faster but are inhibited once the background electrons are Yeated to relativistic temperatures in the shock precursor (Lemoine Pelletier 2011).," The filamentation instability has time to grow only if $Y_{\rm inst}\,\ll\,1$, while other two stream instabilities may grow faster but are inhibited once the background electrons are heated to relativistic temperatures in the shock precursor (Lemoine Pelletier 2011)."
+ For this reason. we consider only the growth of he filamentation instability in the following.," For this reason, we consider only the growth of the filamentation instability in the following."
+" We detine a threshold value Y. such that if 35,4λος micro-instabilities can grow and allow Fermi cycles to develop. as discussed further below. while if luacλος instabilities cannot grow. hence Fermi cycles do not develop."," We define a threshold value $Y_{\rm c}$ such that if $Y_{\rm inst}<Y_{\rm c}$, micro-instabilities can grow and allow Fermi cycles to develop, as discussed further below, while if $Y_{\rm
+ inst}>Y_{\rm c}$, instabilities cannot grow, hence Fermi cycles do not develop."
+ One must expect dina23. in the early stages of the afterglow. since with ay=0 for κr. and ay=1 for pomor.," One must expect $Y_{\rm inst}>Y_{\rm c}$ in the early stages of the afterglow, since with $\alpha_Y=0$ for $r<r_\times$ and $\alpha_Y=-1$ for $r>r_\times$."
+" PIC simulations indicate that £i,,=&,/O.1~1 (Sironi Spitkovsky 2011)."," PIC simulations indicate that $\xi_{\rm b,-1}\equiv
+\xi_{\rm b}/0.1\simeq 1$ (Sironi Spitkovsky 2011)."
+ Early on. as Yiuaλος micro-instabilities are quenched by advection of the plasma through the shock front. hence he magnetic field is everywhere transverse to the shock normal without substantial inhomogeneity on short scales.," Early on, as $Y_{\rm inst}>Y_{\rm c}$, micro-instabilities are quenched by advection of the plasma through the shock front, hence the magnetic field is everywhere transverse to the shock normal without substantial inhomogeneity on short scales."
+" In his ease. Fermi acceleration cannot develop as particles are advected with the magnetic field lines to the far downstream,"," In this case, Fermi acceleration cannot develop as particles are advected with the magnetic field lines to the far downstream."
+ evertheless. the electrons acquire part of the kinetic energy of he incoming protons in the shock transition (as viewed in the shock frame).," Nevertheless, the electrons acquire part of the kinetic energy of the incoming protons in the shock transition (as viewed in the shock frame)."
+ A detailed understanding of this process is still acking but current PIC simulations confirm the above. even in the absence of filamentation in the precursor.," A detailed understanding of this process is still lacking but current PIC simulations confirm the above, even in the absence of filamentation in the precursor."
+" In. particular. Siront Spitkovsky (2011) observe that ο, reaches the value of 0.1 at a magnetization a,=10I. for sa,o20 and larger."," In particular, Sironi Spitkovsky (2011) observe that $\epsilon_e$ reaches the value of $0.1$ at a magnetization $\sigma_{\rm u}=10^{-4}$, for $\gamma_{\rm sh}\simeq 20$ and larger."
+ We adopt this value inthe following., We adopt this value inthe following.
+" For simplicity. we model the shock heated electron. distribution as a restricted. powerlaw with ,=Bruin."," For simplicity, we model the shock heated electron distribution as a restricted powerlaw with $\gamma_{\rm max}=3\gamma_{\rm min}$."
+" The minimal Lorentz factor 5444 is then related to ¢,. through μιGondrysnide. with a;= a normalization prefactor of order unity. |which[1 depends (slightly)| on the modelling of the energy distribution: we adopt s= 2.4. an ad-hoe choice here as well motivated by simplicity (i.e. s will not change once Fermi acceleration becomes effective)."," The minimal Lorentz factor $\gamma_{\rm min}$ is then related to $\epsilon_e$ through $\gamma_{\rm min}= \epsilon_e \gamma_{\rm b} (m_p/m_e)a_{\rm s}$ with $a_{\rm s}=\left[(s-2)/(s-1)\right] \left[1- (\gamma_{\rm
+ max}/\gamma_{\rm min})^{1-s}\right] \left[1- (\gamma_{\rm
+ max}/\gamma_{\rm min})^{2-s}\right]^{-1}$ a normalization prefactor of order unity, which depends (slightly) on the modelling of the energy distribution; we adopt $s=2.4$ , an ad-hoc choice here as well motivated by simplicity (i.e. $s$ will not change once Fermi acceleration becomes effective)."
+" Although the electrons are heated. in the shock transition. the magnetic field is only compressed. so that the magnetic field in the blast frame Di,=Εις."," Although the electrons are heated in the shock transition, the magnetic field is only compressed, so that the magnetic field in the blast frame $B_{\rm b} = 4\gamma_{\rm b}
+B_{\rm w}$."
+ In terms of the conventional parameter ερ describing the fraction of energy carried by the magnetic field in the blast. cg=BF(δπο)Boy.," In terms of the conventional parameter $\epsilon_B$ describing the fraction of energy carried by the magnetic field in the blast, $\epsilon_B=B_{\rm b}^2/(8\pi e_{\rm b}) = 2\sigma_{\rm w}$."
+ As the blast Lorentz factors decreases beyond r.. so does Ya. until τοιY. eventually.," As the blast Lorentz factors decreases beyond $r_\times$, so does $Y_{\rm inst}$, until $Y_{\rm inst}<Y_{\rm c}$ eventually."
+ The filamentation instability now has several to many ¢ folds of growth times before the plasma is advected through the shock front., The filamentation instability now has several to many $e-$ folds of growth times before the plasma is advected through the shock front.
+ This has several consequences of importance., This has several consequences of importance.
+ First of all. the upstream electrons are heated in the micro-turbulence in the shock precursor (Spitkovsky 2008. Lemoine Pelletier 2011) and they therefore reach rough equipartition with the incoming protons after the shock transition. as observed in PIC simulations (Sironi Spitkovsky 2011).," First of all, the upstream electrons are heated in the micro-turbulence in the shock precursor (Spitkovsky 2008, Lemoine Pelletier 2011) and they therefore reach rough equipartition with the incoming protons after the shock transition, as observed in PIC simulations (Sironi Spitkovsky 2011)."
+ This implies ο.2 0.5., This implies $\epsilon_e\simeq 0.5$ .
+ Furthermore. a micro-turbulentmagnetic field is generated on skin depth scales up tocg of a few percents.," Furthermore, a micro-turbulentmagnetic field is generated on skin depth scales up to$\epsilon_B$ of a few percents."
+ We adopt cg=0.05 as a fiducial value in. what follows. in," We adopt $\epsilon_B=0.05$ as a fiducial value in what follows, in"
+3302-101kms! in direction /=234°+11? and b=12°+Telο ,"$330 \pm 101 \kms$ in direction $l =
+234^\circ \pm 11^\circ$ and $b=12^\circ \pm 9^\circ$."
+"The 3500 and 4500 samples both show a dipole in a direction compatible with the Great Attractor region at 300°, b~0°."," The 3500 and 4500 samples both show a dipole in a direction compatible with the Great Attractor region at $l \sim 300^\circ$ , $b \sim 0^\circ$."
+" For the HF the direction shifts to slightly higher b, compatible with a sampleshift in the motion towards the Shapley concentration (whichliesatanaveragedistanceof14000kms'!,Bardellietal.1994) at slightly higher galactic latitude."," For the HF sample the direction shifts to slightly higher $b$, compatible with a shift in the motion towards the Shapley concentration \citep[which lies at an average
+distance of 14000 $\kms$,][]{Bardelli} at slightly higher galactic latitude."
+" Furthermore the amplitude of the dipole decreases, as is expected from the Monte Carlo simulations."," Furthermore the amplitude of the dipole decreases, as is expected from the Monte Carlo simulations."
+" For all the three sub samples we find a relatively large contribution from the quadrupole, showing that the local flow has a significant shear component."," For all the three sub samples we find a relatively large contribution from the quadrupole, showing that the local flow has a significant shear component."
+ The result is consistent in with the expectation from the Monte Carlo simulations., The result is consistent in magnitude with the expectation from the Monte Carlo simulations.
+" magnitudeFrom the figures it can also be seen that there is a change in quadrupole direction with redshift, and that the distributions for the 3500 and 4500 samples are bimodal."," From the figures it can also be seen that there is a change in quadrupole direction with redshift, and that the distributions for the 3500 and 4500 samples are bimodal."
+" This could be because the quadrupole is pointing in different directions at the lower and higher end of the included redshift range of the SNe, and hence for the 3500 sample the bimodality is more apparent."," This could be because the quadrupole is pointing in different directions at the lower and higher end of the included redshift range of the SNe, and hence for the 3500 sample the bimodality is more apparent."
+ Even though we quote formal errors in Table 2 the distribution is highly non-Gaussian and the error bars should be taken as indicative only., Even though we quote formal errors in Table 2 the distribution is highly non-Gaussian and the error bars should be taken as indicative only.
+ The Tonryetal.(2003) data set of 98 SNe was analysed by Hudsonetal.(2004) in order to find the local dipole., The \citet{tonry:2003} data set of 98 SNe was analysed by \citet{hudson} in order to find the local dipole.
+" For the part of the sample with v,«6000kms! they found a dipole of v,=376+81kms! towards |=285°, b=—14°."," For the part of the sample with $v_r < 6000 \kms$ they found a dipole of $v_r = 376 \pm 81\kms$ towards $l = 285^\circ$, $b=-14^\circ$."
+" They do not quote error bars on this result, but for the part of the sample with v,>6000kms! the stated errors are +17° "," They do not quote error bars on this result, but for the part of the sample with $v_r > 6000 \kms$ the stated errors are $\pm 17^\circ$ "
+1,.
+o Brielly. we selected data from the PN detector which satisfies the following criteria: pattern «4 (i.e. single aud double pixel events only). events with pulse heieht between 0.2 and 15 keV. FLAG=0 (the most stringent filler recommended to get a high-quality spectrum).," Briefly, we selected data from the PN detector which satisfies the following criteria: pattern $\leq$ 4 (i.e. single and double pixel events only), events with pulse height between 0.2 and 15 keV, FLAG=0 (the most stringent filter recommended to get a high-quality spectrum)."
+ The red line in Fie., The red line in Fig.
+ 1. shows the 9 point moving averaged light curve extracted from EPIC PN data. from a circular region of 45 arcsecond radius around the source.," \ref{fig:lc} shows the 9 point moving averaged light curve extracted from EPIC PN data, from a circular region of 45 arcsecond radius around the source."
+ A circular. region. of radius 43.5 arcseconds on (he same chip where (he source was located. was chosen to extract background spectra.," A circular, source-free region, of radius 43.5 arcseconds on the same chip where the source was located, was chosen to extract background spectra."
+ We used the SAS task to anlvze the event pattern information near (he source region from the PN event file., We used the SAS task to anlyze the event pattern information near the source region from the PN event file.
+ In the absence of any pile up. Che observed-to-model singles and doubles pattern [Iractions ratios (obtained fromepalplot output) should both be consistent with 1.0 within statistical errors.," In the absence of any pile up, the observed-to-model singles and doubles pattern fractions ratios (obtained from output) should both be consistent with 1.0 within statistical errors."
+ Presence of pile up causes (he singles ratio to decrease [rom unity and the doubles ratio to increase from unity., Presence of pile up causes the singles ratio to decrease from unity and the doubles ratio to increase from unity.
+ When a circular source region is chosen around the f[Iux centroid ofJ10344-396.. the observed distribution of single aud double pattern Iractions (as a Iunction of energv) is hiehlv discrepant Irom the expected model curves.," When a circular source region is chosen around the flux centroid of, the observed distribution of single and double pattern fractions (as a function of energy) is highly discrepant from the expected model curves."
+ The 0.52 keV observed-to-mocdel Iraction of singles in this case is 0.905+0.003 and that of the doubles is 1.298+0.007. showing a hieh pile up.," The 0.5–2 keV observed-to-model fraction of singles in this case is $0.905\pm0.003$ and that of the doubles is $1.298\pm0.007$, showing a high pile up."
+ Therelore we excised the PSF core where pile up is strongest., Therefore we excised the PSF core where pile up is strongest.
+ The optimal racdius-of-exclusion was determined by slowly increasing (he exclusion radius till the observed distribution of single and double events matched fairly well with the expectedcurves?., The optimal radius-of-exclusion was determined by slowly increasing the exclusion radius till the observed distribution of single and double events matched fairly well with the expected.
+. Thus we verified that the source region used by Middletonοἱal.(2009).. ie. excluding; the inner 32 arcsec and using an annular region with outer radius of 45 aresec. is an optimal choice.," Thus we verified that the source region used by \citet{middletonetal2009}, i.e. excluding the inner 32 arcsec and using an annular region with outer radius of 45 arcsec, is an optimal choice."
+ The 0.52 keV observed-to-moclel fraction of singles after the exclusion of inner 32 arcsec is 0.983£0.017 and the corresponding fraction of doubles is 1.0906+0.023., The 0.5–2 keV observed-to-model fraction of singles after the exclusion of inner 32 arcsec is $0.983\pm0.017$ and the corresponding fraction of doubles is $1.096\pm0.028$.
+ The average PN source count-rates in 0.3.10 keV οποιον range before and alter excluding the PSF core were 5.8 counts/s and 0.31 counts/s respectively., The average PN source count-rates in 0.3–10 keV energy range before and after excluding the PSF core were $5.8$ counts/s and $0.31$ counts/s respectively.
+ Appropriate photon redistribution matrices and ancillary region files were created using ancl tasks., Appropriate photon redistribution matrices and ancillary region files were created using and tasks.
+ As noted by Gierliüskiοἱal.(2008) (also see Fig., As noted by \citet{g08} (also see Fig.
+ 1 where the 9-point moving average light curve computed from the 100-s binned PN data is shown by the thick red line). a visual inspection of the Leht curve clearly shows the periodicity of 1 hour.," \ref{fig:lc} where the 9-point moving average light curve computed from the 100-s binned PN data is shown by the thick red line), a visual inspection of the light curve clearly shows the periodicity of $\sim$ 1 hour."
+ These authors also noted that the variability is not strictly periodic over the entire observation. in particular during (he initial ~25 ks of the observation.," These authors also noted that the variability is not strictly periodic over the entire observation, in particular during the initial $\sim$ 25 ks of the observation."
+ Therefore we adopted the folowing algorithm (also illustrated in Fig. 2)), Therefore we adopted the following algorithm (also illustrated in Fig. \ref{fig:gti}) )
+" to extract spectral data from the ""high"" (crest) and ""low (trough) phases.", to extract spectral data from the “high” (crest) and “low” (trough) phases.
+ Let the 9-point moving averaged {lux at time / be denoted by (1)., Let the 9-point moving averaged flux at time $t$ be denoted by $f(t)$ .
+ (MeClintock&Remillard2004 (Irudolvubovοἱal.2001:Osborne2001:Koneetal.2002:Williams2004.. (Williamsetal.2004.2005a.b.d.c).. (Williamsetal.2004.2005a.b.c).. (Williamsetal.2005d)..," \citealp{mcclintock2004} \citealp{trudolyubov2001,osborne2001,kong2002,williams2004hrc}, \citep{williams2004hrc,williams2005bh1,williams2005bh2,williams2005bh4,williams2005bh6}. \citep{williams2004hrc,williams2005bh1,williams2005bh2,williams2005bh6}. \citep{williams2005bh4}."
+ D=24.7. , $B=24.7$ 
+gives an estimate of the contamination factor for this star of 0.015. meaning that less than [1.5 percent of the light in the observing mask ean be from contaminating background stars.,"gives an estimate of the contamination factor for this star of 0.015, meaning that less than 1.5 percent of the light in the observing mask can be from contaminating background stars."
+ This makes it exceedingly improbable that any of the light variations we discuss in the paper might be from a contaminating background star., This makes it exceedingly improbable that any of the light variations we discuss in the paper might be from a contaminating background star.
+ To determine the effective temperature of the star. Balmer lines profiles were compared with synthetic profiles starting with parameters close to those from Kepler Input Catalogue.," To determine the effective temperature of the star, Balmer lines profiles were compared with synthetic profiles starting with parameters close to those from Kepler Input Catalogue."
+ Synthetic spectra were calculated with the program of ?.., Synthetic spectra were calculated with the program of \citet{piskunov92}.
+ We used model atmospheres from the NEMO database (Vienna New Model Grid of Stellar Atmospheres) (2)., We used model atmospheres from the NEMO database (Vienna New Model Grid of Stellar Atmospheres) \citep{heiteretal02}.
+. This grid has an effective temperature step of KK. which we interpolated to get models with a KK step.," This grid has an effective temperature step of K, which we interpolated to get models with a K step."
+ The spectral line list for analysis and calculations was taken from the Vienna Atomic Line Database (VALD. ?)). which includes lines of rare earth elements from the DREAM database ¢?}.," The spectral line list for analysis and calculations was taken from the Vienna Atomic Line Database (VALD, \citealt{kupkaetal99}) ), which includes lines of rare earth elements from the DREAM database \citep{biemontetal99}."
+. We found a best fit for Ha and H.? with synthetic profiles for Zia= T20OKK. with estimated error of KK. The determination of the surface gravity is difficult at this temperature.," We found a best fit for $\alpha$ and $\beta$ with synthetic profiles for $T_{\rm 
+eff} = 7200$ K, with estimated error of K. The determination of the surface gravity is difficult at this temperature."
+" The Balmer line profiles are relatively insensitive to losg. constraining it only to within £0.5,"," The Balmer line profiles are relatively insensitive to $\log g$, constraining it only to within $\pm 0.5$."
+ Ionisation equilibrium for and lines and and lines is potentially much more sensitive for determining logg. but we did not get reliable results with our initial study.," Ionisation equilibrium for and lines and and lines is potentially much more sensitive for determining $\log g$, but we did not get reliable results with our initial study."
+ We expect a more detailed analysis. which is underway. will better constrain logg.," We expect a more detailed analysis, which is underway, will better constrain $\log g$."
+ We therefore use here the photometrically determined value of logg=3.6 (cgs) from the KIC., We therefore use here the photometrically determined value of $\log g = 3.6$ (cgs) from the KIC.
+" Because the spectrum is not very peculiar. we estimate a [c error of 0.3,"," Because the spectrum is not very peculiar, we estimate a $\sigma$ error of $\pm 0.3$."
+ A detailed abundance analysis is in progress., A detailed abundance analysis is in progress.
+ We made preliminary estimates of the abundances of some elements at two rotation phases to get an overview of the spectral character of 110195926., We made preliminary estimates of the abundances of some elements at two rotation phases to get an overview of the spectral character of 10195926.
+ Using synthetic spectra calculated with for model 7;=7200 KK and logg=3.6 we examined some abundances using the first and third NOT spectra listed in refobslog.., Using synthetic spectra calculated with for model $T_{\rm eff} = 7200$ K and $\log g = 3.6$ we examined some abundances using the first and third NOT spectra listed in \\ref{obslog}.
+ The spectral lines are clearly variable with rotation. so other spectra will give values that differ for some elements.," The spectral lines are clearly variable with rotation, so other spectra will give values that differ for some elements."
+ Surprisingly for à roAp star. there are no obvious lines of rare earth elements.," Surprisingly for a roAp star, there are no obvious lines of rare earth elements."
+ The usual strong lines in roAp stars of and are present and indicate overabundances compared to the Sun. but less than is typical for most other roAp stars.," The usual strong lines in roAp stars of and are present and indicate overabundances compared to the Sun, but less than is typical for most other roAp stars."
+ lines are normal., lines are normal.
+ There is no obvious core-wing anomaly (2) in Ha. a usual signature of roAp stars.," There is no obvious core-wing anomaly \citep{cowleyetal01} in $\alpha$, a usual signature of roAp stars."
+ refabund gives the abundances we found compared to solar values., \\ref{abund} gives the abundances we found compared to solar values.
+ The second column gives the abundances for a spectrum obtained at rotational phase 0.31., The second column gives the abundances for a spectrum obtained at rotational phase 0.31.
+ For this spectrum no clear evidence lines of rare earth elements was found., For this spectrum no clear evidence lines of rare earth elements was found.
+ There are some hints of the oresenee of such lines in the spectrum. but the better S/N spectra are required.," There are some hints of the presence of such lines in the spectrum, but the better S/N spectra are required."
+ The fourth column gives abundances for a spectrum obtained at rotational phase 0.95., The fourth column gives abundances for a spectrum obtained at rotational phase 0.95.
+ This spectrum definitely shows ines of and a trace of weak lines of Which confirm the peculiarity of this star., This spectrum definitely shows lines of and a trace of weak lines of which confirm the peculiarity of this star.
+ Lines of strontium are much wider and stronger in this spectrum while the barium lines are much weaker when compared with the spectrum at phase 0.31., Lines of strontium are much wider and stronger in this spectrum while the barium lines are much weaker when compared with the spectrum at phase 0.31.
+ Iron lines are weaker and slightly wider with ower abundances., Iron lines are weaker and slightly wider with lower abundances.
+ The spectral lines show strong variability with rotational period. demonstrating the spotted surface structure.," The spectral lines show strong variability with rotational period, demonstrating the spotted surface structure."
+ By comparing with the roAp star 999563 (2)... we judge that lines of and are formed in rather small spots.," By comparing with the roAp star 99563 \citep{freyhammeretal09}, we judge that lines of and are formed in rather small spots."
+ Other elements also demonstrate nonuniform distribution over stellar surface., Other elements also demonstrate nonuniform distribution over stellar surface.
+ reftig:sp6645 shows the line at four rotation phases., \\ref{fig:sp6645} shows the line at four rotation phases.
+" While better S/N is needed for further study. it is clear that Is concentrated is at least two small spots close to the sulsation pole Gind hence probably close to the magnetic pole) of his star,"," While better S/N is needed for further study, it is clear that is concentrated is at least two small spots close to the pulsation pole (and hence probably close to the magnetic pole) of this star."
+" 110195926 has been observed by the mission during three ""rolls"" up to this writing.", 10195926 has been observed by the mission during three “rolls” up to this writing.
+ Following each quarter of a solar orbit the telescope is rolled to keep its solar panels facing the Sun., Following each quarter of a 370-d solar orbit the telescope is rolled to keep its solar panels facing the Sun.
+ Choices and changes of targets are made on the basis of these “quarters”. which sometimes are split into approximately |-month “thirds”.," Choices and changes of targets are made on the basis of these “quarters”, which sometimes are split into approximately 1-month “thirds”."
+ Most data (for 0000 stars) are obtained in long cadence (LC) with integration times of mmin., Most data (for $>$ 000 stars) are obtained in long cadence (LC) with integration times of min.
+ For 512 targets data are obtained in short cadence (SC), For 512 targets data are obtained in short cadence (SC)
+tecent cliscoveries of quasars at 2cBS (ee. Fan et al.,Recent discoveries of quasars at $z>5.8$ (e.g. Fan et al.
+ 2000. 2001) are finally allowing quantitative studies of the high-redshift InterCalaetic Medium: (16M) ane its reionization jstoryv.," 2000, 2001) are finally allowing quantitative studies of the high-redshift InterGalactic Medium (IGM) and its reionization history."
+ In particular. the detection. of a Gunn-DPeterson rough (Gunn Peterson 1965) in the Ixeck (Decker οἱ al.," In particular, the detection of a Gunn-Peterson trough (Gunn Peterson 1965) in the Keck (Becker et al."
+ 2001) and VET (Penterieci et al., 2001) and VLT (Pentericci et al.
+ 2002) spectra ofthe Sloan Dieital Sky Survey quasar SDSS 1030-0524 at 2=6.28 is a clear indication that the universe is approaching the reionization epoch at z6., 2002) spectra of the Sloan Digital Sky Survey quasar SDSS 1030-0524 at $z=6.28$ is a clear indication that the universe is approaching the reionization epoch at $z\sim 6$.
+ Whatever the exact value of the reionization recdshift. tien db ds Clear that the hydrogen in the ICM is in a highly ionized state at zZ6.," Whatever the exact value of the reionization redshift, $z_{ion}$, it is clear that the hydrogen in the IGM is in a highly ionized state at $z \simlt 6$."
+ The nature of the responsible sources is the subject of a lively debate., The nature of the responsible sources is the subject of a lively debate.
+ Several authors (Macau. llaardt Rees 1999 ancl references therein) have claimed that the known populations of quasars and galaxies provide 10 times fewer ionizing photons than are necessary (o explain the observed. GM ionization level.," Several authors (Madau, Haardt Rees 1999 and references therein) have claimed that the known populations of quasars and galaxies provide $\sim 10$ times fewer ionizing photons than are necessary to explain the observed IGM ionization level."
+ Thus. additional sources of ionizing photons are required at high redshift. the most promising being carly galaxies and quasars.," Thus, additional sources of ionizing photons are required at high redshift, the most promising being early galaxies and quasars."
+ While no evidence for the existence of high-redshift quasars has vet been found. recent observations of temperature ancl metal abundance in the lowest density regions of the LGAL suggest the existence of an carly population of pregalactic stellar objects. which may have contributed to the reionization and metal enrichment of the LGAL (Cowie Songaila 1998: Ellison et al.," While no evidence for the existence of high-redshift quasars has yet been found, recent observations of temperature and metal abundance in the lowest density regions of the IGM suggest the existence of an early population of pregalactic stellar objects, which may have contributed to the reionization and metal enrichment of the IGM (Cowie Songaila 1998; Ellison et al."
+ 2000: Schave ct al., 2000; Schaye et al.
+ 2000: Macau. Ferrara Rees 2001).," 2000; Madau, Ferrara Rees 2001)."
+ For this reason. most theoretical work on IGM reionization has assumed stellar sources.," For this reason, most theoretical work on IGM reionization has assumed stellar sources."
+ The study of. GM. retonization bv. primeval stellar sources has been tackled by several authors. both via semi-analvtic (e.g. Hlaiman Loeb 1997: Valageas Silk 1999: Aliralcla-Escuclé.. Lachnelt Rees 2000: Cojazzi et al.," The study of IGM reionization by primeval stellar sources has been tackled by several authors, both via semi-analytic (e.g. Haiman Loeb 1997; Valageas Silk 1999; Miralda-Escudé,, Haehnelt Rees 2000; Cojazzi et al."
+ 2000) and numerical (c.g. Cinedin Ostriker 1997: Clardi ct al., 2000) and numerical (e.g. Gnedin Ostriker 1997; Ciardi et al.
+ 2000: Chiu Ostriker 2000: Cinedin 2000: Itazoumov οἱ al., 2000; Chiu Ostriker 2000; Gnedin 2000; Razoumov et al.
+ 2002) approaches., 2002) approaches.
+ Two main ingredients are required for a proper treatment of the reionization process: à model of ealaxy formation and2 a reliable treatment of the radiative transfer of ionizing photons., Two main ingredients are required for a proper treatment of the reionization process: a model of galaxy formation and a reliable treatment of the radiative transfer of ionizing photons.
+ The commonly accepted. scenario for structure, The commonly accepted scenario for structure
+We divide our discussion of the results into (hree subsections.,We divide our discussion of the results into three subsections.
+ First. we present the results for the individual galaxies.," First, we present the results for the individual galaxies."
+ Next. we discuss the homogeneity of the structural properties of the galaxies.," Next, we discuss the homogeneity of the structural properties of the galaxies."
+ Finally. we estimate the stellar mass fraction. mass-to-light ratios and the rate ol galaxy. evolution.," Finally, we estimate the stellar mass fraction, mass-to-light ratios and the rate of galaxy evolution."
+" Fieure 1 shows contours for the goodness of fit of the models to the velocity dispersion. measured for each galaxy as a fincetion of the stellar mass fraction f, and the concentration 6 once we have normalized the mass inside the Einstein radius."," Figure \ref{fig:figure1} shows contours for the goodness of fit of the models to the velocity dispersion, measured for each galaxy as a function of the stellar mass fraction $f_*$ and the concentration $c$ once we have normalized the mass inside the Einstein radius."
+ For these caleulations. we have included our estimates of the svstematic errors in the velocity dispersions but used the stated uncertainties in (he Iuminosities.," For these calculations, we have included our estimates of the systematic errors in the velocity dispersions but used the stated uncertainties in the luminosities."
+ Note that the dispersion measurements cannot determine the halo concentrations but the goodness of fit contours always pass through the region set by our prior on (he concentration., Note that the dispersion measurements cannot determine the halo concentrations but the goodness of fit contours always pass through the region set by our prior on the concentration.
+ The permitted stellar mass fractions vary widely between objects., The permitted stellar mass fractions vary widely between objects.
+" Three of the 22 objects. SDSS JOTS7+321. SDSS J12504+052 and PGII15+030. appear (o require mass distributions (hat are more centrally concentrated (han the stars. in the sense that the best fits [or f,<1 have 47>2."," Three of the 22 objects, SDSS $0737+321$, SDSS $1250+052$ and $1115+080$, appear to require mass distributions that are more centrally concentrated than the stars, in the sense that the best fits for $f_* \leq 1$ have $\chi^2>2$."
+ This is also seen in the LSD models for PGI115+080 (Tren&Koopmans2002a).. where the only models consistent with both (he lensing constraint and the estimated velocity dispersion are more centrally concentrated than the stars.," This is also seen in the LSD models for $1115+080$ \citep{tk02a}, where the only models consistent with both the lensing constraint and the estimated velocity dispersion are more centrally concentrated than the stars."
+" A fourth lens. SDSS J1627—005. is only marginally consistent with f,€1."," A fourth lens, SDSS $-$ 005, is only marginally consistent with $f_* \leqslant 1$."
+" Of the remaining 18 galaxies. eleven are consistent with f,=1 GN?< 1). and seven are not."," Of the remaining 18 galaxies, eleven are consistent with $f_* = 1$ $\Delta\chi^2<1$ ), and seven are not."
+ Four of these eleven galaxies have enormous parameter uncertainties., Four of these eleven galaxies have enormous parameter uncertainties.
+ One problem [or manv SLACS lenses is that the scales of the velocity dispersion aperture/elfective radius differ little [rom the observed Einstein radius. which limits the leverage for constraining the mass prolile.," One problem for many SLACS lenses is that the scales of the velocity dispersion aperture/effective radius differ little from the observed Einstein radius, which limits the leverage for constraining the mass profile."
+ Fieuree 3. shows the eeoodness of fit to the mass-to-lighte ratio of each egalaxy egiven the best fit model for the average evolution of the sample., Figure \ref{fig:figure3} shows the goodness of fit to the mass-to-light ratio of each galaxy given the best fit model for the average evolution of the sample.
+ Most of the galaxies are consistent with this best fit model for the mass-to-lisht ratio and its evolution relsec:barma))., Most of the galaxies are consistent with this best fit model for the mass-to-light ratio and its evolution \\ref{sec:barma}) ).
+ The mass-to-lelt ratios of the sample appear to be more uniform than the dynamical properties. probably for the same reasons that there is little scatter in the fundamental plane (see Bernardietal. 2003b)).," The mass-to-light ratios of the sample appear to be more uniform than the dynamical properties, probably for the same reasons that there is little scatter in the fundamental plane (see \citealt{bernardi03b}) )."
+ However. there are three 36 outliers in the sample. SDSS J1420+602. SDSS J1250+052 and 1115434535. all of which have very low M/L ratios compared to the other galaxies.," However, there are three $\sigma$ outliers in the sample, SDSS J1420+602, SDSS J1250+052 and H1543+535, all of which have very low M/L ratios compared to the other galaxies."
+ Note that only one of these. SDSS J1250+052. is also an outlier in the dynamical fits.," Note that only one of these, SDSS J1250+052, is also an outlier in the dynamical fits."
+ This is not unique to our approach. since our mass-to-lieht ratio lor1115434-535 is comparable to that in Treu&Koopmans (2004)..," This is not unique to our approach, since our mass-to-light ratio forH1543+535 is comparable to that in \citet{tk04}. ."
+ In Figure 6 of, In Figure 6 of
+of high-contrast imaging can therefore decisively reject their Sirius C hypothesis.,of high-contrast imaging can therefore decisively reject their Sirius C hypothesis.
+" Although Benest&Duvent(1995) provide no error estimation, we consider a false alarm the most likely explanation for their results."," Although \citet{benest95} provide no error estimation, we consider a false alarm the most likely explanation for their results."
+" Precision astrometry is known to suffer from a multitude of systematic errors ssubtle changes in pixel scale and orientation, and differential atmospheric refraction which depends on airmass, parallactic and ambient conditions), and has led to a series of angle,spurious detections in the past (e.g. Pravdo& citetbeanl0 and Lazorenkoetal.(2011);; (1969) citetgatewood73))."," Precision astrometry is known to suffer from a multitude of systematic errors subtle changes in pixel scale and orientation, and differential atmospheric refraction which depends on airmass, parallactic angle, and ambient conditions), and has led to a series of spurious detections in the past (e.g. \citet{pravdo09}
+ \\citet{bean10} and \citet{lazorenko11}; \citet{vandekamp69} \\citet{gatewood73}) )."
+" Leaving astrometric predictions aside, we can also explore the parameter space for other semi-major axes."," Leaving astrometric predictions aside, we can also explore the parameter space for other semi-major axes."
+" Figure 5 shows the completeness as a function of semi-major axis a = (0.25, 0.50, ..., 10.0} AAU and companion mass m, assuming a flat distribution in eccentricity."," Figure \ref{f:completeness_planets} shows the completeness as a function of semi-major axis $a$ = $\{0.25,$ $0.50,$ $\ldots,$ $10.0\}$ AU and companion mass $m$, assuming a flat distribution in eccentricity."
+" Although our data are sensitive to planets down to an inner working angle of 07 and down to Mj, at large separations, the completeness values drop quickly at shorter separations for a Mjyp object at a = AAU)."," Although our data are sensitive to planets down to an inner working angle of $0\farcs7$ and down to $M_\textrm{Jup}$ at large separations, the completeness values drop quickly at shorter separations for a $M_\textrm{Jup}$ object at $a$ = AU)."
+" This coincides with the domain of long-term stable planet orbits predicted by Holman&Wiegert(1999), with a critical semi-major axis ac= AAU."," This coincides with the domain of long-term stable planet orbits predicted by \citet{holman99}, with a critical semi-major axis $a_\textrm{c}=2.17$ AU."
+" Therefore, plenty of parameter space remains for unseen planets around Sirius A. The upcoming next generation of high-contrast instruments, such as SPHERE (Beuzitetal.2010),, will offer smaller inner working angles and better contrast performance, and thus stand a good chance to detect such planets."," Therefore, plenty of parameter space remains for unseen planets around Sirius A. The upcoming next generation of high-contrast instruments, such as SPHERE \citep{beuzit10}, will offer smaller inner working angles and better contrast performance, and thus stand a good chance to detect such planets."
+" In particular, due to its extreme proximity and brightness, Sirius is the third most promising target (after α Centauri A and B) for the direct detection of exoplanets in reflected light with the SPHERE ZIMPOL imaging (Thalmannetal.2008)."," In particular, due to its extreme proximity and brightness, Sirius is the third most promising target (after $\alpha$ Centauri A and B) for the direct detection of exoplanets in reflected light with the SPHERE ZIMPOL imaging polarimeter \citep{thalmann08}."
+". -1 One thingpolarimeter to keepin mind is the fact that Sirius B was originally a ~5 Mo progenitor star that expanded into a supergiant —125 MMyr ago, with potentially dramatic consequences for the system architecture (Liebertetal.2005)."," -1 One thing to keepin mind is the fact that Sirius B was originally a $\sim$ $M_\odot$ progenitor star that expanded into a supergiant $\sim$ Myr ago, with potentially dramatic consequences for the system architecture \citep{liebert05}."
+". Accretion of ejected material from Sirius B may have caused planets around Sirius A to gain mass and heat, migrate, or form in the first place as second-generation planets (e.g.Perets2010)."," Accretion of ejected material from Sirius B may have caused planets around Sirius A to gain mass and heat, migrate, or form in the first place as second-generation planets \citep[e.g.][]{perets10}."
+". Since these processes would leave the planets hotter and brighter than their unperturbed MMyr- counterparts, our detectable planet mass curves in Figure are conservative for such objects."," Since these processes would leave the planets hotter and brighter than their unperturbed Myr-old counterparts, our detectable planet mass curves in Figure \ref{f:contrast} are conservative for such objects."
+" We present three high-contrast imaging datasets of the Sirius system, all of which reach detection performances in the planetary regime Myyp at1”,, Μι at 2”,, Myyp beyond 4""))."," We present three high-contrast imaging datasets of the Sirius system, all of which reach detection performances in the planetary regime $M_\textrm{Jup}$ at, $M_\textrm{Jup}$ at , $M_\textrm{Jup}$ beyond )."
+ This constitutes an improvement of an order of magnitude in detectable planet mass., This constitutes an improvement of an order of magnitude in detectable planet mass.
+" Taken up to 4.3 years apart, the observations allow us to refute the existence of a substellar companion with a mass of S50 My in a 6.3- orbit as predicted from astrometry measurements of the Sirius AB system (Benest&Duvent1995)."," Taken up to 4.3 years apart, the observations allow us to refute the existence of a substellar companion with a mass of $\lesssim$ $M_\textrm{Jup}$ in a 6.3-year orbit as predicted from astrometry measurements of the Sirius AB system \citep{benest95}."
+". For a companion mass above Myuyp, the chances of a triple false negative at a 5c threshold are0-4%,, on eccentricity."," For a companion mass above $M_\textrm{Jup}$, the chances of a triple false negative at a $\sigma$ threshold are, depending on eccentricity."
+" For the special case of coplanar orbits, the dependingprobability is down to "," For the special case of coplanar orbits, the probability is down to $M_\textrm{Jup}$ )."
+"However, we note that our observations leave open the Mjup).possibility for Jupiter- and Neptune-sized planets around Sirius A, especially at short angular separations."," However, we note that our observations leave open the possibility for Jupiter- and Neptune-sized planets around Sirius A, especially at short angular separations."
+" Furthermore, we confirm the absence dust around Sirius B by the lack of an infrared excess at jum within our precision of mmag."," Furthermore, we confirm the absence dust around Sirius B by the lack of an infrared excess at $\mu$ m within our precision of mag."
+ We thank David Lafreniérre for having provided us with the source code for his LOCI generouslyalgorithm., We thank David Lafrenièrre for generously having provided us with the source code for his LOCI algorithm.
+" We are grateful for the privilege ofobserving from Mauna Kea, which holds great cultural significance for the Hawai*ian indigenous community. AO188),, (Clio).."," We are grateful for the privilege ofobserving from Mauna Kea, which holds great cultural significance for the Hawai`ian indigenous community. , ."
+more interesting. the shock passage of nmultiphase gas with a large density coutrast necessarily vields a multiply curved shock surface that causes vorticity injection ito later accreted gas on the correspouding curvature scales.,"more interesting, the shock passage of multiphase gas with a large density contrast necessarily yields a multiply curved shock surface that causes vorticity injection into later accreted gas on the corresponding curvature scales."
+ Hence. shearing motions are superposed on various scales which Προς amplification of the magnetic Ποια» on these scales.," Hence, shearing motions are superposed on various scales which implies amplification of the magnetic fields on these scales."
+ More sensitive polarized and total intensity observations of NGC 1265 are needed to coufirii our model predictions with radial polarization vectors of the radio torus and a polarization degree of arouud, More sensitive polarized and total intensity observations of NGC 1265 are needed to confirm our model predictions with radial polarization vectors of the radio torus and a polarization degree of around.
+ Tu addition. future observations of different svstenis siuular to NGC 1265 would be beneficial to get a statistical mcasurement of propertics of the accretions shocks and its pre-shock conditions.," In addition, future observations of different systems similar to NGC 1265 would be beneficial to get a statistical measurement of properties of the accretions shocks and its pre-shock conditions."
+ This provides indirect evidence for the existence of the wariu-hot TOAL, This provides indirect evidence for the existence of the warm-hot IGM.
+ Our work shows the poteutial of these kind of serendipitous events as wars to explore the outer fringes of clusters and dynamical features of accretion shocks that are complementary to X-ray observations., Our work shows the potential of these kind of serendipitous events as ways to explore the outer fringes of clusters and dynamical features of accretion shocks that are complementary to X-ray observations.
+ We thank T. A. Eublin. T. Pfrounuer. M. Sun. and an anonvnious referee for helpful comments on this manuscript aud gratefully. acknowledge the ereat atmosphere at the I&avli Iustitute for Theoretical Physics progriun ou Particle Acceleration in Astroplivsical Plasmas. in Santa Barbara (2009 July 26October 3) where this project was initiated.," We thank T. A. lin, T. Pfrommer, M. Sun, and an anonymous referee for helpful comments on this manuscript and gratefully acknowledge the great atmosphere at the Kavli Institute for Theoretical Physics program on Particle Acceleration in Astrophysical Plasmas, in Santa Barbara (2009 July 26–October 3) where this project was initiated."
+ That program was supported im part by the National Science Foundation under erant uo., That program was supported in part by the National Science Foundation under grant no.
+ PIIYO»5-51161., PHY05-51164.
+ ο. eratefully acknowledges flnancial support of the Klaus Tschira Foundation and the National Science and Eneginecring Research Council of Canada., C.P. gratefully acknowledges financial support of the Klaus Tschira Foundation and the National Science and Engineering Research Council of Canada.
+ TAS. was supported in part by NSF eraut ASTOO908668 aud by the University of Minnesota Supercomputing Iustitute., T.W.J. was supported in part by NSF grant AST0908668 and by the University of Minnesota Supercomputing Institute.
+ Although fairly standard. for completeness we show the derivation of the Njet curvature radius that is caused by au external ram pressure wind due to the motion of the galaxy through the IC," Although fairly standard, for completeness we show the derivation of the jet curvature radius that is caused by an external ram pressure wind due to the motion of the galaxy through the ICM."
+ We denote the mass deusity. velocity. aud radius of the jet by pis. alu respectively.," We denote the mass density, velocity, and radius of the jet by $\rho_\rmn{jet}$ , $v_\rmn{jet}$, and $r_\rmn{jet}$, respectively."
+ The two jets coming out of the active galactic core back to back are assunied to be iitially a cjcylinder of rg.leugth Jj., The two jets coming out of the active galactic core back to back are assumed to be initially a cylinder of length $l_\rmn{jet}$.
+ Each jet is then bout over a beudiug radius rà by the rau pressure wind of mass density and velocity ριο and c., Each jet is then bent over a bending radius $r_b $ by the ram pressure wind of mass density and velocity $\rho_\rmn{ICM}$ and $v$.
+" We equate the jet momentuan piosCieRisaha With the transverse force due to the ram pressure wind that acts over a jet propagation timescale (along the bended path in steady state). pxMp2r5lier(2044). to obtain the following equality Solving for the ratio of bending-to-jet radius. we obtain where we introduced the Mach of the jet auc the galaxy. Mj= aud Maa=cfe the adiabatic expoucuts of jet aud surrounding ICAL nunbers744=1/3 aud 51681=5/2. aud their ονpressures,σον P4 aud Pies."," We equate the jet momentum $\rho_\rmn{jet}
+v_\rmn{jet} \pi r_\rmn{jet}^2 l_\rmn{jet}$ with the transverse force due to the ram pressure wind that acts over a jet propagation timescale (along the bended path in steady state), $\rho_\rmn{ICM}v^2 2 r_\rmn{jet} l_\rmn{jet} \pi r_b
+/(2v_\rmn{jet})$, to obtain the following equality Solving for the ratio of bending-to-jet radius, we obtain where we introduced the Mach numbers of the jet and the galaxy, $\M_\rmn{jet}=v_\rmn{jet}/c_\rmn{jet}$ and $\M_\rmn{gal}=v/c_\rmn{ICM}$, the adiabatic exponents of jet and surrounding ICM, $\gamma_\rmn{jet}=4/3$ and $\gamma_\rmn{ICM}=5/3$, and their pressures, $P_\rmn{jet}$ and $P_\rmn{ICM}$."
+ When a low-density bubble crosses a shock of speed ey iu the ICAL the shock accelerates iuto the bubble. pulling yost-shock ambicut gas with it.," When a low-density bubble crosses a shock of speed $v_\rmn{si}$ in the ICM, the shock accelerates into the bubble, pulling post-shock ambient gas with it."
+ The original bubble-ICAL CD follows the shock at a speed intermediate to that of the incident audbubble shocks when theincident shock is at least moderately strong., The original bubble-ICM CD follows the shock at a speed intermediate to that of the incident and bubble shocks when the incident shock is at least moderately strong.
+ Because the shock iutrusiou begins sooner and inpacts more stronely on the leading edge of a roundbubble than its periphery. the aubient eas peuctrates he center of thebubble fist.," Because the shock intrusion begins sooner and impacts more strongly on the leading edge of a round bubble than its periphery, the ambient gas penetrates the center of the bubble first."
+ An initially spheroidalbubble will then evolve iuto a torus (vortex ring) on a timescale determined by the crossing time of the CD through thebubble., An initially spheroidal bubble will then evolve into a torus (vortex ring) on a timescale determined by the crossing time of the CD through the bubble.
+ The shock-produced dynamics in the bubble cau be estimated from the exact solution to the 1D Ricmanu woblem of a plane shock iapacting a CD., The shock-produced dynamics in the bubble can be estimated from the exact solution to the simple 1D Riemann problem of a plane shock impacting a CD.
+" This is found most simply in the initial rest frame of the}sanenibble. which we assunie to be in pressure equilibrium with the wnshocked ICM. and has a density. p,= 6p;. with à<1. propaga"," This is found most simply in the initial rest frame of the bubble, which we assume to be in pressure equilibrium with the unshocked ICM and has a density, $\rho_b =
+\delta \rho_i$ , with $\delta <1$."
+"teAt impact a orward shock will penetrate iuto the low densitybubble with speed. ο, while a rarefaction will backward into the post-shock ICAL"," At impact a forward shock will penetrate into the low density bubble with speed, $v_\rmn{sb}$, while a rarefaction will propagate backward into the post-shock ICM."
+ The bubbleICAL CD will move forward at the same speed as the post-shock flow in the nbble., The bubble-ICM CD will move forward at the same speed as the post-shock flow in the bubble.
+ The full Ricnann solution is obtained by matching the pressure behind the forward shock inside thebubble o the pressure at the foot of the rarefaction iu the ICAL, The full Riemann solution is obtained by matching the pressure behind the forward shock inside the bubble to the pressure at the foot of the rarefaction in the ICM.
+ We define the ICM aud bubble adiabatic indices as 5; and 25 respectively., We define the ICM and bubble adiabatic indices as $\gamma_i$ and $\gamma_b$ respectively.
+ Assuming the shock is propagating from he left. we have right to left four uniform states. 0. 1. 2. 2. where 0 aud 1 are separated by the forward shock. 1 ancl 2 are separated by the CD. while 2 and 3 are separatedby the reverse rarefaction (note that this numbering scheme differs roni the main body of the paper).," Assuming the shock is propagating from the left, we have right to left four uniform states, 0, 1, 2, 3, where 0 and 1 are separated by the forward shock, 1 and 2 are separated by the CD, while 2 and 3 are separated by the reverse rarefaction (note that this numbering scheme differs from the main body of the paper)."
+" The state ""( reepeeseuts the initial couditiousin thebubble. while represents conditions iu the ICAL post-shock flow."," The state “0” represents the initial conditions in the bubble, while “3” represents conditions in the ICM post-shock flow."
+ We set 7?4=Po and ey=οοςCop.," We set $P_1 = P_2 =
+P_*$ and $v_1 = v_2 = v_*=v_\rmn{CD}$."
+ The initialbubble ver(aud ICAL) xessure is 2). while the initialbubble aud ICM souud speeds are related by ej=(yVEITpuοἱ+0).The Mach umiuibers of the external. ICAL shock aud the internal.bubble shock are Mj=ey/e; audMy=cafes—pM; and we introduced the Mach umuber ratio ji5* ," The initial bubble (and ICM) pressure is $P_0$, while the initial bubble and ICM sound speeds are related by $c_b = c_0 =
+\sqrt{\gamma_b P_0/\rho_0} = c_i\sqrt{\gamma_b/ (\gamma_i \delta)}$.The Mach numbers of the external, ICM shock and the internal, bubble shock are $\mathcal{M}_i = v_\rmn{si}/c_i$ and$\mathcal{M}_b = v_\rmn{sb}/c_b \equiv \mu
+\mathcal{M}_i$ and we introduced the Mach number ratio $\mu$ "
+hat is accentuated whenever there is a siguificaut wisinatch between the solar aud model values for Rew.,that is accentuated whenever there is a significant mismatch between the solar and model values for $\rcz$.
+ The best results for intermeciate (late) aceretion leac o (Óócfc) being about two (three) times the value or the CS9s8 SSAL, The best results for intermediate (late) accretion lead to $\dc$ being about two (three) times the value for the GS98 SSM.
+" Interestingly, the best results for intermediate accretion are frou a metal-rich case (panec iun Fieures 7)) aud for late accretion are from a uetal-poor case (panel in Figures 11))"," Interestingly, the best results for intermediate accretion are from a metal-rich case (panel in Figures \ref{fig:hv0a15t10}) ) and for late accretion are from a metal-poor case (panel in Figures \ref{fig:hv0a30t10}) )."
+ The latter xovides an exception to the general rule of correlate and Reg. stemming from the seusitivitv of Zi; in ate (Ocfe)accretion models to the properties of the accretec uaterial.," The latter provides an exception to the general rule of correlated $\dc$ and $\rcz$ , stemming from the sensitivity of $\zini$ in late accretion models to the properties of the accreted material."
+ We discuss this further in 5.3.., We discuss this further in \ref{sec:cmpothers}.
+ For the early accretion scenario. the results displavec iu panel of Figure 3. for metal-rich accretion suggest hat further improvemeut in (39e/e? might be achieve if Mas>0.06MEL. an," For the early accretion scenario, the results displayed in panel of Figure \ref{fig:hv0a5t10} for metal-rich accretion suggest that further improvement in $\dc$ might be achieved if $\mac > 0.06~\msun$."
+dTo test this. we computed inodels with Z4.=0.030 Adie= 0.075. 0.090. 0.110. ane 1.125 AL...," To test this, we computed models with $\zac=0.030$ and $\mac=$ 0.075, 0.090, 0.110, and 0.125 $\msun$."
+ The results for @efe) are given in Figure 13 (top panel).," The results for $\dc$ are given in Figure \ref{fig:z30ext}
+ (top panel)."
+" A mnüiunuunu is found around AL,=1.090 ALL. vieldiug a value for (@e/e} close to that he GSOs SSM (solid black horizoutal line)."," A minimum is found around $\mac=0.090~\msun$ , yielding a value for $\dc$ close to that of the GS98 SSM (solid black horizontal line)."
+" For aceretion nasses greater than A4,=0.090M. (0c/65 begins to rise again."," For accretion masses greater than $\mac=0.090~\msun$, $\dc$ begins to rise again."
+ Iu the bottom panel the sound speed. profile or the AA.=0.000MI. model is shown (red loug-dashed line)., In the bottom panel the sound speed profile for the $\mac=0.090~\msun$ model is shown (red long-dashed line).
+ It is similar to that of the (205 SSM (ervey solid line). except close to the Suus ceuter. where he metallicity and heli abuudauces in the iiodel are oo low.," It is similar to that of the GS98 SSM (grey solid line), except close to the Sun's center, where the metallicity and helium abundances in the model are too low."
+ The ACSSO9 SSME sound speed profile is also shown., The AGSS09 SSM sound speed profile is also shown.
+ For this model Rey=0.716Ιω. in reasonable agreement with the heloscismic value 0.713+0.001R...," For this model $\rcz= 0.716\, \rsun$, in reasonable agreement with the helioseismic value $0.713\pm0.001\, \rsun$."
+ ILowever. following the general relation between Rey aud VS previously described. this model has 34=0.225. further from the helioseisuic value. Ys=0.2[8540.0035. than the ΑΟ SSM. value (Ys=0.2319. the Πιτ of no accretion).," However, following the general relation between $\rcz$ and $\ys$ previously described, this model has $\ys=0.225$, further from the helioseismic value, $\ys=0.2485\pm0.0035$, than the AGSS09 SSM value $\ys=0.2319$, the limit of no accretion)."
+ Therefore. larger accretion masses do not offer a global solution to the solar abundance problem.," Therefore, larger accretion masses do not offer a global solution to the solar abundance problem."
+" All calculations presented in Section 1. have been conrputed assuniue a fixed duration for the accretion phase of Ar,=10 Mw."," All calculations presented in Section \ref{sec:results} have been computed assuming a fixed duration for the accretion phase of $\Delta \tau_{\rm
+ ac}=10$ Myr."
+" Estimated lifetimes of protoplauetary disks are about a few Myr (ILhüschetal.2001:Willams&Cieza2011) so that longer timescales for accretion secum unlikely,"," Estimated lifetimes of protoplanetary disks are about a few Myr \citep{haisch:2001,williams:2011} so that longer timescales for accretion seem unlikely."
+ Ou the other haud. shorter timescales are possible. perhaps even favored.," On the other hand, shorter timescales are possible, perhaps even favored."
+ We therefore repeated the caleulatious from using Ary=1 Alyy. to test the consequences of shorter accretion times.," We therefore repeated the calculations from \ref{sec:results}
+ using $\Delta \tau_{\rm ac}=1$ Myr, to test the consequences of shorter accretion times."
+ As one udelt expect intuitively from an examination of Figure L. the effect of a reduced Ary or given Zac. Mac; aud Το Is analogous to enliauciug he dilution factor of the accreted material: that is. for wo otherwise equivalent accretion scenarios with the sale starting fines. the one with the shorter duration will deposit material iuto a larger mean convection-zone uass diving the accretion. aud thus that material will experience iore dilution.," As one might expect intuitively from an examination of Figure \ref{fig:mconv}, the effect of a reduced $\Delta \tau_{\rm ac}$ for given $\zac$, $\mac$, and $\taui$ is analogous to enhancing the dilution factor of the accreted material: that is, for two otherwise equivalent accretion scenarios with the same starting times, the one with the shorter duration will deposit material into a larger mean convection-zone mass during the accretion, and thus that material will experience more dilution."
+ For example. results for (Tae4. Ara J=(15 Myr.1 Myr) are intermediate between those or (5 Myr.10 Myr. ) aud Abr.dO Myr) shown iu Figures 3. and 7.. respectively.," For example, results for $\taui$, $\Delta \tau_{\rm ac}$ )=(15 Myr,1 Myr) are intermediate between those for (5 Myr,10 Myr ) and (15 Myr,10 Myr), shown in Figures \ref{fig:hv0a5t10} and \ref{fig:hv0a15t10}, respectively."
+ For late accretion. e.9.. after the model has settled ou the MS. the duration of he accretion phase is irrelevant. provided it is short compared to eravitational settling timescales: at late nues. there is little evolution in either the depth or he mass of the convective zone.," For late accretion, e.g., after the model has settled on the MS, the duration of the accretion phase is irrelevant, provided it is short compared to gravitational settling timescales: at late times, there is little evolution in either the depth or the mass of the convective zone."
+ Thus we conclude that changes in the duration of the accretion will not alter our basic conclusion: no combination of the accretion variables Mac. Zac aud Tre will lead to a iiocdoel in which all helioscisumic predictions are improved.," Thus we conclude that changes in the duration of the accretion will not alter our basic conclusion: no combination of the accretion variables $\mac$, $\zac$ and $\taui$ will lead to a model in which all helioseismic predictions are improved."
+ We have compared model solar neutrino fluxes with solar fluxes (column 3 in Table 2))., We have compared model solar neutrino fluxes with solar fluxes (column 3 in Table \ref{tab:neutrinos}) ).
+ The latter have been derived from a combined analysis of all solar neutriuo experiuents with the addition of the solar luuinosity (Borexinocollaboration2011)., The latter have been derived from a combined analysis of all solar neutrino experiments with the addition of the solar luminosity \citep{borex:2011}.
+. The analysis includes the CN-cvcle fluxes. although only upper IHuits are currently available.," The analysis includes the CN-cycle fluxes, although only upper limits are currently available."
+ We fud. using the new nuclear reaction rates from ΕΠ aud the newest Borexino results for the ‘Be flux (Borexinocollaboration2011).. that the CS98 and AGSSO9 SSAIs ave both in excellent aereeimenut with solar neutrino data. producing comparable fits.," We find, using the new nuclear reaction rates from SFII and the newest Borexino results for the $^7$ Be flux \citep{borex:2011}, that the GS98 and AGSS09 SSMs are both in excellent agreement with solar neutrino data, producing comparable fits."
+ These results are summuarized in the last row of Table 2.., These results are summarized in the last row of Table \ref{tab:neutrinos}.
+ Amone the four wellkdetermined fluxes. the D jeutrinos are the most sensitive to accretion. respoudine the variations iu initial core ictallicity that occur j)ecause the accretion imedel is solved under the constraint of a fixed final comvective-zone ictallicity-o-hydrogsen ratio.," Among the four well-determined fluxes, the $^8$ B neutrinos are the most sensitive to accretion, responding to the variations in initial core metallicity that occur because the accretion model is solved under the constraint of a fixed final convective-zone metallicity-to-hydrogen ratio."
+" The ""D solar fux is also the iuost certain. now determined to ~δν "," The $^8$ B solar flux is also the most certain, now determined to $\sim 3\%$."
+Moreover. (0D) xovides Information conrlemeutary to helioscisinologv )ecause the production region. AxzOLR... is essentially inaccessible to p-miode study.," Moreover, $\Phi(^8$ B) provides information complementary to helioseismology because the production region, $R \lesssim 0.1\, \rsun$, is essentially inaccessible to p-mode study."
+ Iu dac we find that all nodels withmetal-rich accretion. because of their lower Zi Values. produce too few *B neutrinos.," In fact, we find that all models withmetal-rich accretion, because of their lower $\zini$ values, produce too few $^8$ B neutrinos."
+" Results are shown in panel of Figures 6.. Εν, and 12.."," Results are shown in panel of Figures \ref{fig:nuv0a5t10}, \ref{fig:nuv0a15t10}, and \ref{fig:nuv0a30t10}."
+ This suggests hat metalvich accretion would create a conflict sinaller han but remiuisceut of the old solar ucutrino problem. ourtieublulv for the imtermeciate and late accretion scenarios.," This suggests that metal-rich accretion would create a conflict smaller than but reminiscent of the old solar neutrino problem, particularly for the intermediate and late accretion scenarios."
+ Iu contrast. models with metal-poor accretion. or an appropriate choice of AZ. that depends on the accretion scenario (carly. intermediate. or late). can be ought iuto excellent agreement with the experimenta fixes.," In contrast, models with metal-poor accretion, for an appropriate choice of $\mac$ that depends on the accretion scenario (early, intermediate, or late), can be brought into excellent agreement with the experimental fluxes."
+" The hieh seusitivitv of the ""D flix to metallicity cads to AM, values as low as 0.01 M. being cisfavorec or the extreme case of late metal-free accretion.", The high sensitivity of the $^8$ B flux to metallicity leads to $\mac$ values as low as 0.01 $\msun$ being disfavored for the extreme case of late metal-free accretion.
+ The “Be. pp. and pep fluxes. are now also determines lugh level of precision. but are less scusitive to 15 core couditious (see Figures 6.. 1O.. aud 12)).," The $^7$ Be, pp, and pep fluxes, are now also determined to a high level of precision, but are less sensitive to the core conditions (see Figures \ref{fig:nuv0a5t10}, \ref{fig:nuv0a15t10}, and \ref{fig:nuv0a30t10}) )."
+ Wit respect to the ΑΟ SSAL a moderate amount οτ uetal-poor accretion cau lead to improvements m beth SB and 'De fluxes. as discussed. below.," With respect to the AGSS09 SSM, a moderate amount of metal-poor accretion can lead to improvements in both the $^8$ B and $^7$ Be fluxes, as discussed below."
+ A slightly nore metal-rich core ποσα» preferred by these fiuxes., A slightly more metal-rich core seems preferred by these fluxes.
+ We have done a global analysis of neutrino fiuxes. witli ιο results of the \? caleulations eiven in panelf of Figures 6.. 10.. and 12..," We have done a global analysis of neutrino fluxes, with the results of the $\chi^2$ calculations given in panel of Figures \ref{fig:nuv0a5t10}, \ref{fig:nuv0a15t10}, and \ref{fig:nuv0a30t10}."
+ In the case of carly accretion. ouly the most metal-rich models with massive accretion show significant disagreement with solar data.," In the case of early accretion, only the most metal-rich models with massive accretion show significant disagreement with solar data."
+ The jeutrino flux responses to accretion are Wore sensitive iu16 intermediate aud late scenarios. however.," The neutrino flux responses to accretion are more sensitive inthe intermediate and late scenarios, however."
+" Figure 11 eives the resulting 47 surfaces iu the M,Zac plane or these scenarios.", Figure \ref{fig:chi2} gives the resulting $\chi^2$ surfaces in the $\mac - \zac$ plane for these scenarios.
+ Iu each panel the best-fit model is denoted by a red diuuoud., In each panel the best-fit model is denoted by a red diamond.
+ Solid lines depict the GS... LOW...and probability contours. and. dotted ines the coutours of probabilityfor the ACSSO9 aud GS9s8 SSAIs )) (Table 2)).," Solid lines depict the , and probability contours, and dotted lines the contours of probabilityfor the AGSS09 and GS98 SSMs ) (Table \ref{tab:neutrinos}) )."
+ In the late accretion case. he top-right corner is ciupty because no solar models satistving our requirements can be constructed for such," In the late accretion case, the top-right corner is empty because no solar models satisfying our requirements can be constructed for such"
+that was used for the MASTERoext analysis.,that was used for the MASTERext analysis.
+ For each siuulatioun we first computed the differcuce between the W- and V-band spectra. and then the mean (inverse noise variance weighted) difference witlin some f-ranec.," For each simulation we first computed the difference between the W- and V-band spectra, and then the mean (inverse noise variance weighted) difference within some $\ell$ -range."
+ Comparing the observed data to these sinulatious. we found that the discrepancy is significant at slightly amore than 30 for both 250<(x600 and LOOx(6€<500. corresponding to the largest reeion of asvuuuetry and to the most discrepant region. respectively.," Comparing the observed data to these simulations, we found that the discrepancy is significant at slightly more than $3\sigma$ for both $250 \le \ell
+\le 600$ and $400 \le \ell \le 500$, corresponding to the largest region of asymmetry and to the most discrepant region, respectively."
+ Prestunably it is possible to boost the sieuificauce by tuniug the region further. but on the other hand. the known I bean amplitude uncertainty that is taken into account bv the WALAP likelihood code will reduce the siguficance by a few tenths of a siena.," Presumably it is possible to boost the significance by tuning the region further, but on the other hand, the known $\sim1\%$ beam amplitude uncertainty that is taken into account by the WMAP likelihood code will reduce the signficance by a few tenths of a sigma."
+ Independent of such nmünor effects. it seenis that the probability of this being a statistical fluke is rather small.," Independent of such minor effects, it seems that the probability of this being a statistical fluke is rather small."
+ After the publication of the prescut paper. a follow-up studv by IEIuffenbereeretal.(2006) revisited the unresolved point source analysis performed— by the team.," After the publication of the present paper, a follow-up study by \citet{huffenberger:2006} revisited the unresolved point source analysis performed by the team."
+ The main result from that work was a best-fit unresolved point source spectrum amplitude of A=0.ULIs?sx (relative to LL GIIz). which is significantly lower than the value of A=0.017j/IKr initially reported CTinshawetal.9006) aud used iu the present paper.," The main result from that work was a best-fit unresolved point source spectrum amplitude of $A =
+0.011 \,\mu\textrm{K}^2 \textrm{sr}$ (relative to 41 GHz), which is significantly lower than the value of $A = 0.017 \,\mu\textrm{K}^2
+\textrm{sr}$ initially reported \citep{hinshaw:2006} and used in the present paper."
+ Thus. a relatively larger contribution is subtracted from the V-baud than from the W-hband spectrum. and in effect. the V-band spectrum has been over-corrected.," Thus, a relatively larger contribution is subtracted from the V-band than from the W-band spectrum, and in effect, the V-band spectrum has been over-corrected."
+ After taking into account this new wuplitude. the discrepancy between the two baud spectra was found to be siguificant at ~26 using the same test as above.," After taking into account this new amplitude, the discrepancy between the two band spectra was found to be significant at $\sim2\sigma$ using the same test as above."
+ While the quoted point source correction can account for a substantial amount of the observed discrepancy. there ids still a small difference present. and this could possibly indicate further systematics.," While the quoted point source correction can account for a substantial amount of the observed discrepancy, there is still a small difference present, and this could possibly indicate further systematics."
+ In this respect. if could be worth cousideriug unnmodceled beam axvuuuctrics.," In this respect, it could be worth considering unmodeled beam asymmetries."
+ As a preliminary step in the three-vear analysis. Tinshawetal.(2006) considered the impact of such asviunuetries on the individual cross-spectra.," As a preliminary step in the three-year analysis, \citet{hinshaw:2006} considered the impact of such asymmetries on the individual cross-spectra."
+ After an extensive analysis. they concluded that their magnitude is less than at (6x1000. and therefore sufficiently sinall to neglect in further analvses.," After an extensive analysis, they concluded that their magnitude is less than at $\ell \le 1000$, and therefore sufficiently small to neglect in further analyses."
+ However. at the relevant scales. the absolute amplitude of the temperature power spectrum is- about 2000-2gs. aud a beam axviuuetry bias of therefore corresponds to an expected discrepancy of 20pl?," However, at the relevant scales, the absolute amplitude of the temperature power spectrum is about $2000\,\mu\textrm{K}^2$, and a beam asymmetry bias of therefore corresponds to an expected discrepancy of $20\,\mu\textrm{K}^2$."
+ This corresponds roughly to lo of the V- versus W-band difference. aud. if determined appropriate. its correction could bring the overall difference well within the statistical errors.," This corresponds roughly to $1\sigma$ of the V- versus W-band difference, and, if determined appropriate, its correction could bring the overall difference well within the statistical errors."
+ Fortunately. this issue will be clearer with additional vears of observations.," Fortunately, this issue will be clearer with additional years of observations."
+ Tn the previous two sections we cousicdered the cluperature angular power spectrum as observed byWAZAP.. and our main conclusion from these analyscs is that the speetruii is reasonable at all angular scales.," In the previous two sections we considered the temperature angular power spectrum as observed by, and our main conclusion from these analyses is that the spectrum is reasonable at all angular scales."
+ Tlowever. there are small but clearly noticeable vases at both laree and small scales.," However, there are small but clearly noticeable biases at both large and small scales."
+ In this section. we seek to quantify the iupact of this bias in terms of he cosmological paraineters for a ual six-paraieter ACDM model.," In this section, we seek to quantify the impact of this bias in terms of the cosmological parameters for a minimal six-parameter $\Lambda$ CDM model."
+" The combined effect of both the low- aud Hel-f discrepancies are studied by IHuffeubergeretal. (2006).. aud the effect on extended cosimological models (ο,ον, models inchiding massive neutrinos and rine of he spectral iudex) are considered by I&xxistiausenctal.(2006)."," The combined effect of both the low- and $\ell$ discrepancies are studied by \citet{huffenberger:2006}, and the effect on extended cosmological models (e.g., models including massive neutrinos and running of the spectral index) are considered by \citet{kristiansen:2006}."
+. We perform four sets of similar analyses. all primarily sed ou the likelihood code.," We perform four sets of similar analyses, all primarily based on the likelihood code."
+ First. we adopt the likelihood. as provided.," First, we adopt the likelihood as provided."
+ Second. we replace the low-f ikchhood (both the pixel-based estimator aud the low-f ATASTER estimator) with a BlackwellRao (BR) Cibbs-sed estimator for (6<30 (Chuctal.2005).," Second, we replace the $\ell$ likelihood (both the pixel-based estimator and the $\ell$ MASTER estimator) with a Blackwell-Rao (BR) Gibbs-based estimator for $\ell \le 30$ \citep{chu:2005}."
+. Third. we do the same for (<12 alone.," Third, we do the same for $\ell \le 12$ alone."
+ Finally. we use the Mido=16. μας=30 pixel-based likelihood for the V-wand and Wp2e cut described earlier.," Finally, we use the $N_{\textrm{side}} = 16$, $\ell_{\textrm{max}} = 30$ pixel-based likelihood for the V-band and Kp2e cut described earlier."
+ For all cases. we analyze two data combinations: the data alone. aud the combination ofWALAP.. Achar (I&uoetal.2001) and BOOAIERanG (Moutrovetal.2005:Piaceutinial.2005:Joneset 2005).," For all cases, we analyze two data combinations; the data alone, and the combination of, Acbar \citep{kuo:2004} and BOOMERanG \citep{montroy:2005,piacentini:2005,jones:2005}."
+. Mareiualization of SZ was rot imcluded., Marginalization of SZ was not included.
+ Note that we use the likelihood at high Cs in all cases in order to liehlieht the effects of the ow-f likelihood bias., Note that we use the likelihood at high $\ell$ 's in all cases in order to highlight the effects of the $\ell$ likelihood bias.
+ While using a Cubbs sampling sed estimator even at hieh Cs could potentially have veneficial effects in terms of uncertainties. it might confuse the low-f bias analysis by introducing other differences.," While using a Gibbs sampling based estimator even at high $\ell$ 's could potentially have beneficial effects in terms of uncertainties, it might confuse the $\ell$ bias analysis by introducing other differences."
+ The cosmological pariuneters corresponding to these ikelihoods were established through standard Markov Chain Monte Carlo analysis., The cosmological parameters corresponding to these likelihoods were established through standard Markov Chain Monte Carlo analysis.
+" Iun keeping with our shilosophy of cross-verification. two independent codes were used for ai few cases,"," In keeping with our philosophy of cross-verification, two independent codes were used for a few cases."
+ In the first case. CosimoMC. (Lewis&Bridle2002) was dowloaded aud appropriately modified. aud in the second. a stand-alone code was written from scratch in €||.," In the first case, CosmoMC \citep{lewis:2002} was downloaded and appropriately modified, and in the second, a stand-alone code was written from scratch in C++."
+ The two codes returned identical distributions. aud as usual we show only one set of results in the following.," The two codes returned identical distributions, and as usual we show only one set of results in the following."
+ We also performed simular analyses using oulv tenmiperature-data. imposing a Gaussian prior ou the optical depth of 7=0.1040.03 to simulate the effect of polarization data. but without relying on the accuracy of these;," We also performed similar analyses using only temperature-data, imposing a Gaussian prior on the optical depth of $\tau = 0.10\pm0.03$ to simulate the effect of polarization data, but without relying on the accuracy of these."
+ As expected. we then," As expected, we then"
+As incutioned in the last section. it seenis desirable to use ouly forbidden lines for studies of the σας metallicity of NLRs in Sls.,"As mentioned in the last section, it seems desirable to use only forbidden lines for studies of the gas metallicity of NLRs in S1s."
+ Furthermore. it seems ideal to use only low-ionization forbidden cussion lines to discuss the gas properties.," Furthermore, it seems ideal to use only low-ionization forbidden emission lines to discuss the gas properties."
+" This is because low-ionization line-cinitting regions and hiel-iouization line-ciuitting regions niav be spatially scerceated (c.g.. Baker 1997: Muraviuuna Taniguchi 1998a. 1998b: Nagao. Tanuiguehi. Muraviuna 2000: Nagao. Muraviuua. Taniguchi 2001b. 2001c: sec also Tes. Barthel. Foshbury 1993),"," This is because low-ionization line-emitting regions and high-ionization line-emitting regions may be spatially segregated (e.g., Baker 1997; Murayama Taniguchi 1998a, 1998b; Nagao, Taniguchi, Murayama 2000; Nagao, Murayama, Taniguchi 2001b, 2001c; see also Hes, Barthel, Fosbury 1993)."
+" Takine the above two constraints into account. we examine whether or not the forbidden cmission-line flux ratios of |O I[AG300/[N. λύσσα, [ο n]A3727/[N. u|A6583.. aud. 15 H[AAGTIT.GT21/|N 11]AG583 can be useful estimators for the uitrogen abundance of gas clouds in the NLRs."," Taking the above two constraints into account, we examine whether or not the forbidden emission-line flux ratios of [O $\lambda$ 6300/[N $\lambda$ 6583, [O $\lambda$ 3727/[N $\lambda$ 6583, and [S $\lambda \lambda$ 6717,6731/[N $\lambda$ 6583 can be useful estimators for the nitrogen abundance of gas clouds in the NLRs."
+ Iu Figure 5. we show the relationship among the three forbidden emissiou-Iine flux ratios.," In Figure 5, we show the relationship among the three forbidden emission-line flux ratios."
+ Hore we use the data of Nagao et al. (, Here we use the data of Nagao et al. (
+20016) in order to fud possible correlatious among the three fux ratios.,2001c) in order to find possible correlations among the three flux ratios.
+ Since the data are not corrected for the dust reddening. we also show the effect of the reddening correction. which is calculated by adopting the extinction curve of Cardelli et al. (," Since the data are not corrected for the dust reddening, we also show the effect of the reddening correction, which is calculated by adopting the extinction curve of Cardelli et al. ("
+1989). in Figure 5.,"1989), in Figure 5."
+ As shown apparently. positive correlations are secu among the three dus ratios.," As shown apparently, positive correlations are seen among the three flux ratios."
+ The corresponding correlation coefficients are 0.67 Land 0.561 for [O Πλ τον AG583 versus ο Πλ n|AG583. aud |S. n]AAGT1T.6731/[N II[AG5823 versus [ο YAG300/[N AG5823. respectively.," The corresponding correlation coefficients are 0.674 and 0.561 for [O $\lambda$ 3727/[N $\lambda$ 6583 versus [O $\lambda$ 6300/[N $\lambda$ 6583 and [S $\lambda \lambda$ 6717,6731/[N $\lambda$ 6583 versus [O $\lambda$ 6300/[N $\lambda$ 6583, respectively."
+ These correlation cocficicuts Πρίν that there are mneoaniusful correlations aniong the three flux ratios., These correlation coefficients imply that there are meaningful correlations among the three flux ratios.
+ What makes these correlatious?, What makes these correlations?
+ Or. can these correlatious be interpreted as sequences of the uitrogen abuudance?," Or, can these correlations be interpreted as sequences of the nitrogen abundance?"
+" Im order to investigate these issues, we perform photoionization model calculations."," In order to investigate these issues, we perform photoionization model calculations."
+ The model methods and the results are preseuted iu the following sections., The model methods and the results are presented in the following sections.
+ To investigate the origin of the correlations presented i the last section. we carry out photoionization ποσο] oealeulatious by using the publicly available code Cloudy version 91.00 (Ferland 1997. 2000).," To investigate the origin of the correlations presented in the last section, we carry out photoionization model calculations by using the publicly available code $Cloudy$ version 94.00 (Ferland 1997, 2000)."
+ Παο we assume uifori density eas clouds with a plane-parallel eeoimetzy., Here we assume uniform density gas clouds with a plane-parallel geometry.
+ The parameters for the calculations are (I) the hydrogen cusity of a cloud (ayy). (IH) the ionization parameter (E). TT) the chemical composition of the eas. aud (IV) the VAhape of the spectral euergv distribution (SED) of the oeiput coutinuun radiation.," The parameters for the calculations are (I) the hydrogen density of a cloud $n_{\rm H}$ ), (II) the ionization parameter $U$ ), (III) the chemical composition of the gas, and (IV) the shape of the spectral energy distribution (SED) of the input continuum radiation."
+ We perform several model ruus covering the following ranges of parameters: 10279 αιP x pg x 10%? cm? aud LO35 « DU < 20., We perform several model runs covering the following ranges of parameters: $^{2.0}$ $^{-3}$ $\leq$ $n_{\rm H}$ $\leq$ $^{5.0}$ $^{-3}$ and $^{-3.5}$ $\leq$ $U$ $\leq$ $^{-2.0}$.
+ For the chemical composition of the gas clouds. we asstune the case that the metals are all scaled keeping solar proportions except for nitrogen: the nitrogen. abundance scales with Z? (this assimuption is altered in some calculations presented i 85).," For the chemical composition of the gas clouds, we assume the case that the metals are all scaled keeping solar proportions except for nitrogen; the nitrogen abundance scales with $Z^2$ (this assumption is altered in some calculations presented in 5)."
+" We calculate the models covering a inetalliitv range of 0.25 < ZZ. sx 3.0, which corresponds to 8.27 < 12 | log (O/II) € 9.35 and 6.77 < 12 | log (NAD) € 8.92."," We calculate the models covering a metallicity range of 0.25 $\leq$ $Z/Z_{\odot}$ $\leq$ 3.0, which corresponds to 8.27 $\leq$ 12 + log (O/H) $\leq$ 9.35 and 6.77 $\leq$ 12 + log (N/H) $\leq$ 8.92."
+ The adopted elemental abuudances of the solar ones relative to hydrogen are taken from Crevesse Anders (1989) with exteusions, The adopted elemental abundances of the solar ones relative to hydrogen are taken from Grevesse Anders (1989) with extensions
+"direction using all the optical photo-z members, and thus, is a surface galaxy density.","direction using all the optical photo-z members, and thus, is a surface galaxy density."
+" We separate LFs using similar criteria, logustn>2 (dense), 1.6<logXsu,2 (intermediate), and logXsi,«1.6 (sparse), then plot LFs of each region in the stars, circles, and squares in Fig.2.."," We separate LFs using similar criteria, $log \Sigma_{5th}\geq 2 $ (dense), $1.6\leq log \Sigma_{5th}<2$ (intermediate), and $log \Sigma_{5th}< 1.6$ (sparse), then plot LFs of each region in the stars, circles, and squares in \ref{fig:8umlf_density}."
+ A fraction of the total volume of the cluster is assigned to each density group in inversely proportional to the sum of Y of each group., A fraction of the total volume of the cluster is assigned to each density group in inversely proportional to the sum of $\Sigma_{5th}^{3/2}$ of each group.
+" Interestingly, the faint-end slope becomes flatter and flatter with decreasing local galaxy density."," Interestingly, the faint-end slope becomes flatter and flatter with decreasing local galaxy density."
+ This result is consistent with our comparison with the field in Fig.l.., This result is consistent with our comparison with the field in \ref{fig:8umlf}.
+" In fact, the lowest density LF (squares) has a flat faint-end tail similar to that of the field LF."," In fact, the lowest density LF (squares) has a flat faint-end tail similar to that of the field LF."
+" Since these LFs are based on the same data, changes in the faint-end slope are not likely due to the errors in completeness correction nor calibration problems."," Since these LFs are based on the same data, changes in the faint-end slope are not likely due to the errors in completeness correction nor calibration problems."
+ The completeness of the deep and shallow regions of the cluster are measured separately., The completeness of the deep and shallow regions of the cluster are measured separately.
+ The changes in the slope is much larger than the maximum completeness correction of25%., The changes in the slope is much larger than the maximum completeness correction of.
+". We also checked the cluster LFs as a function of cluster centric radius, to find no significant difference, perhaps due to the elongated morphology of this cluster."," We also checked the cluster LFs as a function of cluster centric radius, to find no significant difference, perhaps due to the elongated morphology of this cluster."
+" At the same time, assuming the same cluster volume, Fig.2 shows that a possible contamination from the field galaxies to cluster LFs is only ~0.1% in the dense region and ~1% even in the sparse region."," At the same time, assuming the same cluster volume, \ref{fig:8umlf_density} shows that a possible contamination from the field galaxies to cluster LFs is only $\sim$ in the dense region and $\sim$ even in the sparse region."
+" It is interesting that not just the change in the scale of the LFs, but there is a change in the L* and the faint-end slope (o) of the LFs, resulting in the deficit in the 10.255;<logLay,10.8L for cluster LFs."," It is interesting that not just the change in the scale of the LFs, but there is a change in the $L^*$ and the faint-end slope $\alpha$ ) of the LFs, resulting in the deficit in the $10.2L_{\odot}<\log L_{8\mu m}<10.8L_{\odot}$ for cluster LFs."
+ One might imagine a change just in L* might explain the difference in Fig.1.., One might imagine a change just in $L^*$ might explain the difference in \ref{fig:8umlf}.
+" However, in Fig.2,, there clearly is a change in the slope as a function of X5;;."," However, in \ref{fig:8umlf_density}, there clearly is a change in the slope as a function of $\Sigma_{5th}$."
+" However, interpretation is rather complicated; a shape of LF would not change if field galaxies infall into cluster uniformly without changing their star-formation activity."," However, interpretation is rather complicated; a shape of LF would not change if field galaxies infall into cluster uniformly without changing their star-formation activity."
+" Although in cluster environment, a fraction of MIR luminous galaxies is smaller than field (Koyamaetal.,2008),, uniform and instant quenching of star-formation activity of field galaxies can only shift a LF, but cannot account for a change in L* and a of the LFs."," Although in cluster environment, a fraction of MIR luminous galaxies is smaller than field \citep{2008MNRAS.391.1758K}, uniform and instant quenching of star-formation activity of field galaxies can only shift a LF, but cannot account for a change in $L^*$ and $\alpha$ of the LFs."
+ Two important findings in this work are; (i) L* is smaller in the cluster. (, Two important findings in this work are; (i) $L^*$ is smaller in the cluster. (
+ii) the faint-end slopes become steeper toward higher-density regions.,ii) the faint-end slopes become steeper toward higher-density regions.
+" To explain these changes in LFs, IR-luminous galaxies need to be preferentially reduced, with a relative increase of IR-faint galaxies."," To explain these changes in LFs, IR-luminous galaxies need to be preferentially reduced, with a relative increase of IR-faint galaxies."
+" However, an environmental-driven physical process such as the pressure stripping or galaxy-merging would quench star-formation not only in massive galaxies but in less massive galaxies as well, and thus is not able to explain the observed changes in LFs."," However, an environmental-driven physical process such as the ram-pressure stripping or galaxy-merging would quench star-formation not only in massive galaxies but in less massive galaxies as well, and thus is not able to explain the observed changes in LFs."
+" On the other hand, it has been frequently observed that more massive galaxies formed earlier in the Universe."," On the other hand, it has been frequently observed that more massive galaxies formed earlier in the Universe."
+" This downsizing scenario also depends on the environment, in the sense that galaxies with same mass are more evolved in higher density environments than galaxis in less dense environments (Gotoetal.,2005;Tanaka2005,2008).."," This downsizing scenario also depends on the environment, in the sense that galaxies with same mass are more evolved in higher density environments than galaxis in less dense environments \citep{2005ApJ...621..188G,2005MNRAS.362..268T,2008A&A...489..571T}."
+" Statistically, a good correlation has been found between Lr7R and stellar mass (Elbazetal.,2007).."," Statistically, a good correlation has been found between $L_{TIR}$ and stellar mass \citep{2007A&A...468...33E}."
+" Our finding of the relative lack of IR-Iuminous galaxies in the cluster environment may be consistent with the downsizing scenario, where higher density regions have more evolved galaxies and lacks massive star-forming galaxies."," Our finding of the relative lack of IR-luminous galaxies in the cluster environment may be consistent with the downsizing scenario, where higher density regions have more evolved galaxies and lacks massive star-forming galaxies."
+" In contrast, in lower density regions more massive galaxies are still star-forming."," In contrast, in lower density regions more massive galaxies are still star-forming."
+" However, since the data we have shown is in IR luminosity, to conclude on this, we need good stellar mass estimate based on deeper near-IR data."," However, since the data we have shown is in IR luminosity, to conclude on this, we need good stellar mass estimate based on deeper near-IR data."
+" Although a specific mechanism is unclear, the steep faint- could also result from the enhanced star-formation in less massive galaxies."," Although a specific mechanism is unclear, the steep faint-end could also result from the enhanced star-formation in less massive galaxies."
+" In the above scenario, massive galaxies have already ceased their star-formation in the cluster, but less massive galaxies are still forming stars."," In the above scenario, massive galaxies have already ceased their star-formation in the cluster, but less massive galaxies are still forming stars."
+ These less massive galaxies may stop star-formation soon to join the faint-end of the red-sequence (Koy, These less massive galaxies may stop star-formation soon to join the faint-end of the red-sequence \citep{2007MNRAS.382.1719K}. .
+"amaetal., 2007).. Fig.1 Caputietal.(2007); Bav", \ref{fig:8umlf} \citet{2007ApJ...660...97C}; \citet{2008A&A...479...83B}.
+ouzetetal.(2008).. Caputietal.(2007) (see," \citet{2007ApJ...660...97C} \citep[see Fig.12 of ][]{2007ApJ...660...97C,GOTO_NEP_LF}."
+wo samples being the same is for Dununositv class IIL for luminosity class IV. aud for Dpuuinositv class V. There is therefore no statistical evidence for any das in the esiny values for all three luminosity classes in the sample.,"two samples being the same is for luminosity class III, for luminosity class IV, and for luminosity class V. There is therefore no statistical evidence for any bias in the $v \sin i$ values for all three luminosity classes in the sample."
+ What is clear from Figure 8 is that there is a considerablv lower mean cesiu/ for bhuunmositv class, What is clear from Figure 8 is that there is a considerably lower mean $v \sin i$ for luminosity class
+"the temperature falls with height, upwardly displaced fluid elements are warmer than their new surroundings; they expand and continue to rise.","the temperature falls with height, upwardly displaced fluid elements are warmer than their new surroundings; they expand and continue to rise."
+" Similarly, downwardly displaced fluid elements sink."," Similarly, downwardly displaced fluid elements sink."
+" An order of magnitude calculation shows that displacements grow exponentially, with the growth rate (or e-folding rate) p~|gOlnT/8z|- 3."," An order of magnitude calculation shows that displacements grow exponentially, with the growth rate (or e-folding rate) $p \sim |g \, \partial \ln T / \partial z|^{-1/2}$ ."
+ ? investigated the limit in which the temperature of the plasma increases with height., \citet{Quataert2008} investigated the limit in which the temperature of the plasma increases with height.
+" The instability in this case is known as the heat-flux-driven buoyancy instability, or HBI."," The instability in this case is known as the heat-flux-driven buoyancy instability, or HBI."
+ We sketch the growth of the HBI in the bottom row of Figure 1.., We sketch the growth of the HBI in the bottom row of Figure \ref{fig:buoyancy_schematic}.
+" Here, the initial equilibrium state has vertical magnetic field lines and a constant heat flux; the latter is required for the plasma to be in thermal equilibrium."," Here, the initial equilibrium state has vertical magnetic field lines and a constant heat flux; the latter is required for the plasma to be in thermal equilibrium."
+" However, perturbations to the magnetic field divert the heat flux and conductively heat or cool pockets of the plasma."," However, perturbations to the magnetic field divert the heat flux and conductively heat or cool pockets of the plasma."
+" If the temperature increases with height, upwardly displaced fluid elements become warmer than their surroundings and therefore experience a destabilizing buoyant response."," If the temperature increases with height, upwardly displaced fluid elements become warmer than their surroundings and therefore experience a destabilizing buoyant response."
+ These perturbations grow with the same growth rate InT/0z|2.," These perturbations grow with the same growth rate $|g \, \partial \ln T /
+ \partial z|^{-1/2}$."
+" ?|g performed a WKB analysis on equations in the Boussinesq limit and obtained the following dispersion relation for plasmain a constant gravitational field g=—g2, threaded by a constant magnetic field with any orientation in the ὢ—Z plane: Here, p=—iw is the local growth rate of the mode (the time dependence of the perturbation is e?*), wa=k-va is the Alfvénn crossing frequency, describes the buoyant response of an adiabatic plasma, k.=k/k is the direction of the wave vector, and is inversely proportional to the conduction time across the wavelength of the mode."," \citet{Quataert2008} performed a WKB analysis on equations in the Boussinesq limit and obtained the following dispersion relation for plasmain a constant gravitational field $\vec{g} = -g \, \hat{z}$, threaded by a constant magnetic field with any orientation in the $\hat{x}-\hat{z}$ plane: Here, $p = -i \omega$ is the local growth rate of the mode (the time dependence of the perturbation is $e^{p t}$ ), $\omega_{\mr{A}} =
+ \vec{k}\cdot\vec{v}_{\mr{A}}$ is the Alfvénn crossing frequency, describes the buoyant response of an adiabatic plasma, $\hat{k} = \vec{k}/k$ is the direction of the wave vector, and is inversely proportional to the conduction time across the wavelength of the mode."
+ It is convenient to identify Gbuoy=|gOlnT/8z|V? as a characteristic frequency for the buoyancy instabilities and we will also use thuoy=Wruoy and laa—N! in our analysis.," It is convenient to identify $\omega_{\mr{buoy}} = | g \; \partial \ln
+ T / \partial z |^{1/2}$ as a characteristic frequency for the buoyancy instabilities and we will also use $t_{\mr{buoy}} = \omega_{\mr{buoy}}^{-1}$ and $t_{\mr{ad}} = N^{-1}$ in our analysis."
+" In the limit that conduction is rapid compared to any dynamical response (w,>wa,NN, Wpuoy), equation simplifies As we stated in section 2.1,, the growth rate in the rapid conduction limit is independent of the thermal conductivity."," In the limit that conduction is rapid compared to any dynamical response $\omega_{\kappa} \gg \omega_{\mr{A}}, \; N, \; \omega_{\mr{buoy}}$ ), equation simplifies As we stated in section \ref{subsec:equations}, the growth rate in the rapid conduction limit is independent of the thermal conductivity."
+" Equation shows that magnetic tension can suppress the MTI and HBI; if wi> p? must be negative, and the plasma is stable to small Whuoysperturbations."," Equation shows that magnetic tension can suppress the MTI and HBI; if $\omega_{\mr{A}}^2 > \omega_{\mr{buoy}}^2$, $p^2$ must be negative, and the plasma is stable to small perturbations."
+" In this paper, we focus on the relatively weak field limit in which magnetic tension does not suppress the instabilities, and we take wA<Wbuoy in equation (7).."," In this paper, we focus on the relatively weak field limit in which magnetic tension does not suppress the instabilities, and we take $\omega_{\mr{A}} \ll
+ \omega_{\mr{buoy}}$ in equation ."
+ The dominant role of the magnetic field in our analysis is to enforce anisotropicelectron heat transport., The dominant role of the magnetic field in our analysis is to enforce anisotropicelectron heat transport.
+"For the non-radiative simulations, Το shows the expected behavior: The temperature scales ccL2, and thus T,, as well.","For the non-radiative simulations, $T_c$ shows the expected behavior: The temperature scales $\propto L^2$, and thus $T_c$, as well."
+" The density does not depend on L, its profile is uniquely c1?/? (see?).."," The density does not depend on $L$, its profile is uniquely $\propto r^{-2/3}$ \citep[see][]{Shandarin89}."
+" However, because we increase the number of grid points in order to keep the spatial resolution fixed, larger simulations resolve the central peak better."," However, because we increase the number of grid points in order to keep the spatial resolution fixed, larger simulations resolve the central peak better."
+" In the result, we get an apparent dependence of the central density, i.e., of the density at the innermost grid point, ος12/3."," In the result, we get an apparent dependence of the central density, i.e., of the density at the innermost grid point, $\propto L^{2/3}$."
+" However, the latter reflects only the degree of resolution."," However, the latter reflects only the degree of resolution."
+" If including cooling and heating processes, the above relations are no longer valid."," If including cooling and heating processes, the above relations are no longer valid."
+" The central temperature stays roughly constant at about 2x10 K, weakly decreasing at higher L proportional to 1;035."," The central temperature stays roughly constant at about $2\times10^4$ K, weakly decreasing at higher $L$ proportional to $L^{-0.38}$."
+ The central hydrogen number density shows a strong dependence on L approaching the asymptote ny.οL?*8.," The central hydrogen number density shows a strong dependence on $L$ approaching the asymptote $n_{\mathsc{H}c} \propto
+L^{2.38}$."
+" At the edge of the isothermal core, the temperature strongly increases outwards, but then flattens again."," At the edge of the isothermal core, the temperature strongly increases outwards, but then flattens again."
+" We identify the core radius A, as the distance from the center to the inflexion point of the temperature profile.", We identify the core radius $\lambda_\circ$ as the distance from the center to the inflexion point of the temperature profile.
+" Then the core radius shows a dependence on the scale length proportional to L~!38,", Then the core radius shows a dependence on the scale length proportional to $L^{-1.38}$.
+" The above mentioned scaling relations for Τ., ΠΗς, and ο can be explained using simple thermodynamical arguments."," The above mentioned scaling relations for $T_c$, $n_{\mathsc{H}c}$, and $\lambda_\circ$ can be explained using simple thermodynamical arguments."
+" Actually, the size of the core and its density is fully determined by conditions of hydrostatic equilibrium."," Actually, the size of the core and its density is fully determined by conditions of hydrostatic equilibrium."
+" In this case Euler’s equations yield for the central pressure p.~nycd and Poisson's equation yields ¢~ny,A2."," In this case Euler's equations yield for the central pressure $p_c
+\sim n_{\mathsc{H}c} \phi$ and Poisson's equation yields $\phi \sim n_{\mathsc{H}c}
+\lambda_\circ^2$."
+" Combining those two estimates we obtain Without heating the actual pressure of the cooling gas is not able to withstand the gravity forces, and no distinguishable core is forming in these simulations."," Combining those two estimates we obtain Without heating the actual pressure of the cooling gas is not able to withstand the gravity forces, and no distinguishable core is forming in these simulations."
+" Including heating, the pressure is raised by several orders of magnitude during reionization at6."," Including heating, the pressure is raised by several orders of magnitude during reionization at."
+ In Fig., In Fig.
+" 7 we show the path with respect to the central point of the configuration in a phase space diagram, where the temperature T; is plotted against the hydrogen number density nye (these are physical, not super-comoving quantities)."," \ref{fPhase} we show the path with respect to the central point of the configuration in a phase space diagram, where the temperature $T_c$ is plotted against the hydrogen number density $n_{\mathsc{H}c}$ (these are physical, not super-comoving quantities)."
+ The curve enters the shown domain at the center of the bottom, The curve enters the shown domain at the center of the bottom
+as cau be verified by combining Eqs. b.5.. a,"as can be verified by combining Eqs. \ref{Psi}, , \ref{HopfEquation}, ,"
+ud 6. to reproduce Eq., and \ref{linearOp1} to reproduce Eq.
+ 1. in the absence of any stochastic lorcine., \ref{algebraicEOMs} in the absence of any stochastic forcing.
+ As Eq., As Eq.
+ 5. is linear in V. the averageDm V.=Ov[q(/).\pl) obevs the same equation: however V eucapsulates information about the equal-time moments. as can be seen by repeated differentiation ol Eq.," \ref{HopfEquation} is linear in $\Psi$, the average $\overline{\Psi} \equiv \langle \Psi[q(t),~ p] \rangle$ obeys the same equation; however $\overline{\Psi}$ encapsulates information about the equal-time moments, as can be seen by repeated differentiation of Eq."
+ { with respect to p;. Followed by averaging: (For the case of time-averagiug. the statistics do not vary in time. DV=(. and the statistics are obtained from the solution of the time-independent equation HW=0.)," \ref{Psi} with respect to $p_i$, followed by averaging: (For the case of time-averaging, the statistics do not vary in time, $\frac{\partial}{\partial t} \overline{\Psi} = 0$, and the statistics are obtained from the solution of the time-independent equation $\hat{H} \overline{\Psi} = 0$ .)"
+ The Hopf fuuctional Ψ ay also expressed as the expouential of a power series in p;. the coelficieuts being the eumulants: as can be checked by use of Eq.," The Hopf functional $\overline{\Psi}$ may also expressed as the exponential of a power series in $p_i$ , the coefficients being the cumulants: as can be checked by use of Eq."
+ 7 to reproduce the moments iu terms of the eumulants. Eqs. 3..," \ref{Hopf-moments} to reproduce the moments in terms of the cumulants, Eqs. \ref{cumulants}. ."
+ Stochastic forcing cau now be iucluded with the addition of the D;j term: Upou substituting Eq., Stochastic forcing can now be included with the addition of the $\Gamma_{ij}$ term: Upon substituting Eq.
+ 8 into Eq., \ref{Hopf-cumulants} into Eq.
+ 5 aud collecting powers of p; one obtains the equations of motion (EON) for the cumulauts that truncated at third order read where we defer discussion of the parameter j& until later., \ref{HopfEquation} and collecting powers of $p_i$ one obtains the equations of motion (EOM) for the cumulants that truncated at third order read where we defer discussion of the parameter $\mu$ until later.
+ Here for compactness we lave introduced the short-hand notation {} to denote syminetrization over indices that maintaius synunetries 6j;= aud similarly for the third eunulant., Here for compactness we have introduced the short-hand notation $\{ \}$ to denote symmetrization over indices that maintains symmetries $\dot{c}_{ij} = \dot{c}_{ji}$ and similarly for the third cumulant.
+ Truncated at secoud-order (CE2) the cumulaut expansion is realizable (Salinou1995) and. iu the seuse that tlie euergy density is positive aud the second. cumulant obeys positivity constraints., Truncated at second-order (CE2) the cumulant expansion is realizable \citep{salmon98} and well-behaved in the sense that the energy density is positive and the second cumulant obeys positivity constraints.
+ Coine to third-order (CE3) and. beyond. introduces dilliculties., Going to third-order (CE3) and beyond introduces difficulties.
+ A phenomenological eddy-cdaimping parameter (Orszag1977:AudreL971) ji that moclels the neglect of the fourth aud higher cumulauts from the hierarchy is included in the last of Eq.," A phenomenological eddy-damping parameter \citep{Orszag:1977,Andre:1974p398} $\mu$ that models the neglect of the fourth and higher cumulants from the hierarchy is included in the last of Eq."
+ 10 and is required to prevent blow-up., \ref{ce-eom} and is required to prevent blow-up.
+ This ad-hoc procedure is somewhat unsatisfactory auc more robust inethods may be unecessary., This ad-hoc procedure is somewhat unsatisfactory and more robust methods may be necessary.
+ Iudeeddeterminingreliable methods of trumcatingthe hierarchy.is a matter of current research., Indeeddeterminingreliable methods of truncatingthe hierarchyis a matter of current research.
+"at lo confidence, unless otherwise stated.","at $\sigma$ confidence, unless otherwise stated."
+ All the models include Galactic absorption by a column density NG”10?? cm?(?)..," All the models include Galactic absorption by a column density $N_{\rm
+H}^{\rm{Gal}}=4.35\times 10^{20}$ $^{-2}$."
+ We focused our analysis on the 0.2-10 keV band of the pn data and on the 0.3-8 keV band of the MOS data., We focused our analysis on the 0.2-10 keV band of the pn data and on the 0.3-8 keV band of the MOS data.
+" We simultaneously fit the average spectra of the four epochs, Dec. 2004, Jun. 2008, Dec. 2008, and 2009 Long Look, to investigate the causes of the observed variability on month-long time scales."," We simultaneously fit the average spectra of the four epochs, Dec. 2004, Jun. 2008, Dec. 2008, and 2009 Long Look, to investigate the causes of the observed variability on month-long time scales."
+" For each epoch, the model parameters were kept fixed between the MOS and the pn datasets, except for the primary power-law normalizations that were left free to vary within to account for the instrumental cross-calibration uncertainties."," For each epoch, the model parameters were kept fixed between the MOS and the pn datasets, except for the primary power-law normalizations that were left free to vary within to account for the instrumental cross-calibration uncertainties."
+" For the sake of clarity, in most of the following plot we only show the pn data."," For the sake of clarity, in most of the following plot we only show the pn data."
+" The spectra were initially fit to a simple power-law model, with fixed slope and free amplitude between different epochs."," 	 The spectra were initially fit to a simple power-law model, with fixed slope and free amplitude between different epochs."
+ Results are shown in Fig. 3.., Results are shown in Fig. \ref{FIG3}.
+" The fit is very poor, with a chi square per degrees of freedom y*/v=8801/1591, and a very flat photon index I'~0.6."," The fit is very poor, with a chi square per degrees of freedom $\chi^2/\nu=8801/1591$, and a very flat photon index $\Gamma\sim 0.6$."
+" The spectral residuals of all the four epochs have a similar shape, and indeed allowing for the spectral slope of the power-law to vary between epochs does not improve the fit statistics."," The spectral residuals of all the four epochs have a similar shape, and indeed allowing for the spectral slope of the power-law to vary between epochs does not improve the fit statistics."
+ Three main deviations from the, Three main deviations from the
+ollowing these variations in detail.,following these variations in detail.
+ The sampling was subsequently increased al 14.5 GlIz. and these data reveal a well-defined. monotonic swing in EVPA at 14.5 GllIz during very ale 199Learly 1995 associated wilh the total aud polarized flux event which occurred during this period.," The sampling was subsequently increased at 14.5 GHz, and these data reveal a well-defined, monotonic swing in EVPA at 14.5 GHz during very late 1994-early 1995 associated with the total and polarized flux event which occurred during this period."
+ This result is consistent with the VLBI results of (1996). who find a strong jet feature in their 1995.29 map al 5 GlIz which was much weaker in a LOSO image., This result is consistent with the VLBI results of \citet{pre96} who find a strong jet feature in their 1995.29 map at 5 GHz which was much weaker in a 1989 image.
+ We conclude that variability is not restricted (ο flat spectrum sources and that its characteristic behavior in an individual source can change on time scales of order a decade., We conclude that variability is not restricted to flat spectrum sources and that its characteristic behavior in an individual source can change on time scales of order a decade.
+ Such changes in behavior have previously been identified in UMRAO observations of a few non-sample members. e.g. 3C 120. and OJ 237 (Aller.Aller.andIIughes1996).. both flat spectrum. ecore-dominated AGNs.," Such changes in behavior have previously been identified in UMRAO observations of a few non-sample members, e.g. 3C 120, and OJ 287 \citep{all96}, both flat spectrum, core-dominated AGNs."
+ Multi-epoehi VLBI observations demonstrate that changes in these objects are consistent with the development and evolution of hotspots in the flows of QSO and BL Lac jets., Multi-epoch VLBI observations demonstrate that changes in these objects are consistent with the development and evolution of hotspots in the flows of QSO and BL Lac jets.
+ Variability in galaxies is most often not outburst-like. as is commonly found in DL Lacs and QSOs. but rather consists of slower variations. often only identiliable with a time window of order a lew decades. or inlrequent events. also requiring a long time window to detect. variations in the integrated total flix data.," Variability in galaxies is most often not outburst-like, as is commonly found in BL Lacs and QSOs, but rather consists of slower variations, often only identifiable with a time window of order a few decades, or infrequent events, also requiring a long time window to detect variations in the integrated total flux data."
+ We attribute the behavior we see in galaxies {ο relatively long timescales combined with the low Iuminositv of the active core components relative to that from the extended structure. requiring an unusually bright event for detection.," We attribute the behavior we see in galaxies to relatively long timescales combined with the low luminosity of the active core components relative to that from the extended structure, requiring an unusually bright event for detection."
+ Long term linear polarization observations provide information on (he characteristics of the quiescent [low - namely the degree of turbulence aud (he dominant orientation of the magnetic field in the radio-emitiing region since. in the optically thin case. the magnetic field orientation is orthogonal to the directlv-measiured EVPA.," Long term linear polarization observations provide information on the characteristics of the quiescent flow - namely the degree of turbulence and the dominant orientation of the magnetic field in the radio-emitting region since, in the optically thin case, the magnetic field orientation is orthogonal to the directly-measured EVPA."
+" Detailed modeling, of selected lime segments of our multilrequencey total [lux density and linear polarization observations in three sources exhibiting the expected signature of a propagating shock has provided strong evidence for the passage of a transverse shock in these selected cases 1991).", Detailed modeling of selected time segments of our multifrequency total flux density and linear polarization observations in three sources exhibiting the expected signature of a propagating shock has provided strong evidence for the passage of a transverse shock in these selected cases \citep{hug91}.
+. In other objects. comparison of changes in the EVPA during active phases and the VLBI parsec-seale structural axis. a measure of the flow direction in the jet. has demonstrated that shocks are most often at oblique augles to the flow direction rather (han transverse (ο il.," In other objects, comparison of changes in the EVPA during active phases and the VLBI parsec-scale structural axis, a measure of the flow direction in the jet, has demonstrated that shocks are most often at oblique angles to the flow direction rather than transverse to it."
+ Such comparisons have used both integrated polarization observations. when," Such comparisons have used both integrated polarization observations, when"
+reduction to (H-nbK)~1.340.2 by 2010.,reduction to $\sim$ $\pm$ 0.2 by 2010.
+ Unfortunately. lack of contemporaneous data did not permit us to determine if the outbursts peaked at the same time in both bands.," Unfortunately, lack of contemporaneous data did not permit us to determine if the outbursts peaked at the same time in both bands."
+ Similar. photometric behaviour - a transient outburst characterised by an increase in IR continuum excess with wavelength - is observed in other WC binary systems where dust production rates are modulated on the orbital period and peak during periaston passage (e.g. Williams et al. 1994.. 2001..," Similar photometric behaviour - a transient outburst characterised by an increase in IR continuum excess with wavelength - is observed in other WC binary systems where dust production rates are modulated on the orbital period and peak during periaston passage (e.g. Williams et al. \cite{williams94}, \cite{williams01},"
+ van der Hucht et al. 2001))., van der Hucht et al. \cite{vdh01}) ).
+ Consequently. we infer the presence of an additional long period tertiary companion to W239. with the lack of a second outburst suggesting a period of at least ~6 yr for this component.," Consequently, we infer the presence of an additional long period tertiary companion to W239, with the lack of a second outburst suggesting a period of at least $\sim$ 6 yr for this component."
+ However. the broadband near-IR colours are clearly indicative of the continued presence of hot dust (e.g. Crowther et al. 2006a)):," However, the broadband near-IR colours are clearly indicative of the continued presence of hot dust (e.g. Crowther et al. \cite{pac06}) );"
+ in this respect W239 bears a close resemblance to WR48a. which also shows both persistent and episodic. enhanced dust emission.," in this respect W239 bears a close resemblance to WR48a, which also shows both persistent and episodic, enhanced dust emission."
+ Williams et al. (2003)), Williams et al. \cite{williams03}) )
+ also interpret this as evidence for a triple system. with a short period binary (P «1 yr) component responsible for continous dust formation and an interaction with a tertiary companion in a wide. eccentric orbit (P4 730 yr: P. Williams. priv.," also interpret this as evidence for a triple system, with a short period binary $_{\rm orb}<$ 1 yr) component responsible for continous dust formation and an interaction with a tertiary companion in a wide, eccentric orbit $_{\rm orb}>$ 30 yr; P. Williams, priv."
+ comm., comm.
+ 2010) leading to the outburst in 1980., 2010) leading to the outburst in 1980.
+ We return to this in Sect., We return to this in Sect.
+ 3., 3.
+ However. in this regard we highlight that despite the sparse sampling. the decay in flux from the outburst peak in the nBK band lightcurve appeared to occur more rapidly than that of the the H band.," However, in this regard we highlight that despite the sparse sampling, the decay in flux from the outburst peak in the nBK band lightcurve appeared to occur more rapidly than that of the the H band."
+ In other periodic/episodie dust producing systems the opposite is observed. with more extended decays at progressively longer wavelenths due to the cooling of the dust as it is carried from the system by the stellar winds.," In other periodic/episodic dust producing systems the opposite is observed, with more extended decays at progressively longer wavelenths due to the cooling of the dust as it is carried from the system by the stellar winds."
+ Unfortuntely. given the limitations of our current dataset in particular the lack of longer wavelength data - we cannot presently explain this apparent difference in behaviour.," Unfortuntely, given the limitations of our current dataset - in particular the lack of longer wavelength data - we cannot presently explain this apparent difference in behaviour."
+ Finally. optical observations by Bonanos (2007)) between 2006 June l5-July 25 revealed no evidence for eclipses or elipsoidal. modulation 1n. the lighteurve.," Finally, optical observations by Bonanos \cite{bonanos}) ) between 2006 June 15-July 25 revealed no evidence for eclipses or elipsoidal modulation in the lightcurve."
+ We repeated the experiment with a higher cadence - ~ 1500 individual I band observations between 2008 July 8-17 - but again found no evidence for modulation on the orbital period although. às with Bonanos. intra-night variability of Al< 0.08mag.," We repeated the experiment with a higher cadence - $\sim$ 1500 individual I band observations between 2008 July 8-17 - but again found no evidence for modulation on the orbital period although, as with Bonanos, intra-night variability of $\Delta$ $\leq$ 0.08mag."
+ over short ( hr) timescales was observed., over short $\sim$ hr) timescales was observed.
+ In order to investigate the nature of the WC9 primary of W239 we have employed the non-LTE model atmosphere code CMFGEN (Hillier Miller 1998.. 1999)).," In order to investigate the nature of the WC9 primary of W239 we have employed the non-LTE model atmosphere code CMFGEN (Hillier Miller \cite{hillier98}, \cite{hillier99}) )."
+ To accomplish this. in addition to the data described in Sect.," To accomplish this, in addition to the data described in Sect."
+ 2.1 we also employed the spectra presented by Clark et al. (2002)), 2.1 we also employed the spectra presented by Clark et al. \cite{clark02}) )
+ and Crowther et al. (2006a)), and Crowther et al. \cite{pac06}) )
+ and we refer the reader to those works for a deseription of the observational and reduction. procedures employed., and we refer the reader to those works for a description of the observational and reduction procedures employed.
+ While these data span a significant period of time. no secular spectral evolution is apparent (Sect.," While these data span a significant period of time, no secular spectral evolution is apparent (Sect."
+ 2.1) and there is an excellent correspondence between regions of the spectrum that have been multiply sampled., 2.1) and there is an excellent correspondence between regions of the spectrum that have been multiply sampled.
+ The methodology adopted follows that of the analysis of the WC9 star WRIO3 by Crowther et al. (2006b))., The methodology adopted follows that of the analysis of the WC9 star WR103 by Crowther et al. \cite{wc9}) ).
+ Given the significant reddening to W239 some of the UV/optical transitions utilised by Crowther et al., Given the significant reddening to W239 some of the UV/optical transitions utilised by Crowther et al.
+ were unavailable to us: thus our primary diagnostics were CILlt 6678. 7230. 9900A. Cutt 6740. 8500. 9700. 21100 and rv;t 14300 which. when combined with local continuum levels enabled an estimation of temperature. wind properties and stellar luminosity.," were unavailable to us; thus our primary diagnostics were $\lambda\lambda$ 6678, 7230, 9900, $\lambda\lambda$ 6740, 8500, 9700, 21100 and $\lambda$ 14300 which, when combined with local continuum levels enabled an estimation of temperature, wind properties and stellar luminosity."
+ One complication was the significant dilution of the spectrum at long (>| jim) wavelengths by emission from hot dust. as well as a potential further contribution from à massive companion(s).," One complication was the significant dilution of the spectrum at long $>1$ $\mu$ m) wavelengths by emission from hot dust, as well as a potential further contribution from a massive companion(s)."
+ The former was accounted for by the inclusion of a simple greybody model for the dust emission. assuming a single temperature of 1300 K. To represent the latter we utilised à Kuruez LTE model with T~35 kK - corresponding to an O7 V star - and loge ~4.," The former was accounted for by the inclusion of a simple greybody model for the dust emission, assuming a single temperature of 1300 K. To represent the latter we utilised a Kurucz LTE model with $\sim$ 35 kK - corresponding to an O7 V star - and $g \sim$ 4."
+ We adopted an equal light ratio at 8500 A for the WR and OB companion(s)., We adopted an equal light ratio at 8500 $\AA$ for the WR and OB companion(s).
+ The relative contribution of the O7 V star was determined from the analysis of Martins, The relative contribution of the O7 V star was determined from the analysis of Martins
+between bulegc-dominatec galaxies with too low fractions of available cold gas to ignite central starbursts. sometimes calledm,"between bulge-dominated galaxies with too low fractions of available cold gas to ignite central starbursts, sometimes called."
+ergers. appears that such a transition is necessary to explain properties0t like e.g. the isophotal shape of massive elliptical galaxies (IxXhochfar&Burkert2005)., It appears that such a transition is necessary to explain properties like e.g. the isophotal shape of massive elliptical galaxies \citep{kb05}.
+. To stress the importance of taking into account the past mereing history of progenitors. we investigate in Fie.," To stress the importance of taking into account the past merging history of progenitors, we investigate in Fig."
+ 5. the question of the relative number of mergers between galaxies without bulges to the number of mergers between galaxies with bulges. depending on the masses of the remnants.," \ref{fig2} the question of the relative number of mergers between galaxies without bulges to the number of mergers between galaxies with bulges, depending on the masses of the remnants."
+ As can be seen. it becomes very unlikely to find mergers between two pure cise galaxies that leacl to very massive remnants.," As can be seen, it becomes very unlikely to find mergers between two pure disc galaxies that lead to very massive remnants."
+" At masses above AL,23.LO"" M.i the majority of mergers take place between galaxies that have already. experienced a major merger in their past."," At masses above $M_* \simeq 3 \times 10^{9}$ $_{\odot},$ the majority of mergers take place between galaxies that have already experienced a major merger in their past."
+" Lt is interesting to note that we do not find any mergers between bulec-less disc galaxies in our simulation at masses larger than the characteristic mass scale Ader. suggesting that the elect of constant surface mass density. ff, in elliptical galaxies more massive than Ale: is closely related to the bulge fractions in the progenitor galaxies."," It is interesting to note that we do not find any mergers between bulge-less disc galaxies in our simulation at masses larger than the characteristic mass scale $M_C$, suggesting that the effect of constant surface mass density $\mu_*$ in elliptical galaxies more massive than $M_C$ is closely related to the bulge fractions in the progenitor galaxies."
+ Inspecting the environmental dependence. we find that the ratios evolve independently of the environment of the galaxy and that bulee-less mergers stop occurring at lower masses in dense environments like groups and clusters.," Inspecting the environmental dependence, we find that the ratios evolve independently of the environment of the galaxy and that bulge-less mergers stop occurring at lower masses in dense environments like groups and clusters."
+ This is à consequence of the higher merger fractions at earlier times in these environments (Ixhochfar& 2001).., This is a consequence of the higher merger fractions at earlier times in these environments \citep{kb01}.
+ The populations of stars present in a spheroid at a given time in its evolution are a composite of stars formed in various progenitors at dillerent. times and in cillerent nmioces., The populations of stars present in a spheroid at a given time in its evolution are a composite of stars formed in various progenitors at different times and in different modes.
+ Here we distinguish between two moces. thecomponent” and the dise mode.," Here we distinguish between two modes, the and the disc mode."
+component. he merger component is a composite of all stars that were formed. because of the consumption of cold eas during major mergers in progenitor galaxies., The merger component is a composite of all stars that were formed because of the consumption of cold gas during major mergers in progenitor galaxies.
+ On the other hand. the quiescent component is the amount of stars formed in gaseous dises according to the Schmidt-Ixennicutt law curing the evolution of the progenitor galaxies.," On the other hand, the quiescent component is the amount of stars formed in gaseous discs according to the Schmidt-Kennicutt law during the evolution of the progenitor galaxies."
+ When calculating cach component for a spheroid. we sum up the mass of the merger component so far and subtract it from the spheroid mass to get the quiescent component.," When calculating each component for a spheroid, we sum up the mass of the merger component so far and subtract it from the spheroid mass to get the quiescent component."
+" To investigate svstematic trends in. the fraction of merecr-induced ancl quiescent stellar components. in elliptical galaxies ancl bulges. we use in the following the conditional distribution where AL, is the stellar mass in the spheroid previously formecl quicscenth. A,,; the mass of the spheroid and A; the total stellar mass of the galaxy."," To investigate systematic trends in the fraction of merger-induced and quiescent stellar components in elliptical galaxies and bulges, we use in the following the conditional distribution where $M_q$ is the stellar mass in the spheroid previously formed quiescently, $M_{bul}$ the mass of the spheroid and $M_*$ the total stellar mass of the galaxy."
+" Phe merger component in the spheroid is then simply given by 1AL,Mp.", The merger component in the spheroid is then simply given by $1-M_{q}/M_{bul}$.
+ AM not stated otherwise. this distribution is derived. for the simulated galaxy population at redshifts zx0.3.," If not stated otherwise, this distribution is derived for the simulated galaxy population at redshifts $z\leq 0.3$."
+ The quiescent component of stars in spheroids is getting more dominant with galaxy mass up to masses around Ale when it starts becoming a constant fraction of the stars in the spheroid (Fig. 6))., The quiescent component of stars in spheroids is getting more dominant with galaxy mass up to masses around $M_C$ when it starts becoming a constant fraction of the stars in the spheroid (Fig. \ref{fig3}) ).
+ The behaviour at. galaxy masses below Me: can be understood. by acknowledging that massive spheroids form late in the hierarchical galaxy formation paradigm. thereby allowing more stars to be formecl in progenitor clises between individual major merger events. and that feedback from supernovae is less ellicient in reheating cold gas in discs which are embedded in massive," The behaviour at galaxy masses below $M_C$ can be understood by acknowledging that massive spheroids form late in the hierarchical galaxy formation paradigm, thereby allowing more stars to be formed in progenitor discs between individual major merger events, and that feedback from supernovae is less efficient in reheating cold gas in discs which are embedded in massive"
+indices or the ionization parameters of the absorbers to vary produced: acceptable values of ντἐν for the 5 objects. incdiviclually. ancl as a sample.,"indices or the ionization parameters of the absorbers to vary produced acceptable values of $\chi^{2}/\nu$ for the 5 objects individually, and as a sample."
+ Since the two models are equivalent in terms of X7/6. we can only discriminate between the two models by looking to sce which provides the more physically plausible ancl selfconsistent explanation for the spectra.," Since the two models are equivalent in terms of $\chi^{2}/\nu$, we can only discriminate between the two models by looking to see which provides the more physically plausible and self-consistent explanation for the spectra."
+ In. particular. we can test the hypothesis that. the underlving spectra. of these objects are reasonably tvpical for QSOs. and their unusual observed. spectral shapes are due to absorption.," In particular, we can test the hypothesis that the underlying spectra of these objects are reasonably typical for QSOs, and their unusual observed spectral shapes are due to absorption."
+ Two properties in particular are important in this context: the spectral index of the underlving X-ray spectrum. and the ratio of optical to X-ray flux. which we parameterise as the equivalent eox. which is the energy. index of the power law that would connect the restframe aand 2 keV [lux densities.," Two properties in particular are important in this context: the spectral index of the underlying X-ray spectrum, and the ratio of optical to X-ray flux, which we parameterise as the equivalent $\alpha_{OX}$, which is the energy index of the power law that would connect the restframe and 2 keV flux densities."
+ leginning with the X-ray spectral slope. the most appropriate comparison sample of normal QSOs is the sample examined by Mateosetal.(2005).. which is an ray selected sample with a similar Dux limit to the hhard spectrum. sample from which our targets were drawn. and measured in the same energy range. and with the same instrumentation ELPIC) as our targets.," Beginning with the X-ray spectral slope, the most appropriate comparison sample of normal QSOs is the sample examined by \citet{mateos05a}, which is an X-ray selected sample with a similar flux limit to the hard spectrum sample from which our targets were drawn, and measured in the same energy range, and with the same instrumentation EPIC) as our targets."
+" Mateosetal.(2005). found. that the distribution of QSO spectral slopes can be described as à Gaussian with a mean of ay=0.98 and a standard deviation σι,=0.21.", \citet{mateos05a} found that the distribution of QSO spectral slopes can be described as a Gaussian with a mean of $\alpha_{X}=0.98$ and a standard deviation $\sigma_{\alpha}=0.21$.
+ From Table 3. we see that when the spectral slopes are allowed to vary. the best fit values of ay for aancl aare relatively close to the mean of the ax cistribution.," From Table \ref{tab:modelfits} we see that when the spectral slopes are allowed to vary, the best fit values of $\alpha_{X}$ for and are relatively close to the mean of the $\alpha_{X}$ distribution."
+ llowever. the best ft values of ἂν lorINJ094144..1921509.. and. aare all unusually low. Iving within the lowest 3 percent of the cüstribution.," However, the best fit values of $\alpha_{X}$ for, and are all unusually low, lying within the lowest 3 percent of the distribution."
+ This is an unlikely circumstance if the intrinsic photon indices of our sample were drawn from the same distribution as QSOs in general. and suggests that the cold. photoelectrie absorption does not. Lully account for the unusual. X-ray spectral shapes observed in. these objects.," This is an unlikely circumstance if the intrinsic photon indices of our sample were drawn from the same distribution as QSOs in general, and suggests that the cold photoelectric absorption does not fully account for the unusual X-ray spectral shapes observed in these objects."
+ In comparison. the spectral slopes are fixed in the ionized. absorber model. and so are by design consistent with a typical underlving QSO spectrum.," In comparison, the spectral slopes are fixed in the ionized absorber model, and so are by design consistent with a typical underlying QSO spectrum."
+ Aloving to a@ox. we have caleulated this quantity for each of the model fits ancl each of the QSOs.," Moving to $\alpha_{OX}$, we have calculated this quantity for each of the model fits and each of the QSOs."
+ In each case the restframe flux has been estimated. from the D magnitude listed in Page.Mittaz&Carrera(2001).. assuming an optical slope of ao=0.5.," In each case the restframe flux has been estimated from the B magnitude listed in \citet{page01b}, assuming an optical slope of $\alpha_{O}=0.5$."
+ The 2 keV Bux is determined. from the continuum model of the X-ray spectral fit. ie. it is the intrinsic continuum Lux. corrected for the modelled. X-ray absorption.," The 2 keV flux is determined from the continuum model of the X-ray spectral fit, i.e. it is the intrinsic continuum flux, corrected for the modelled X-ray absorption."
+ The distribution of aox in normal QSOs is known to depend on Luminosity (Vignali.Brandt&Scehneider.2003:Stratevaetal..2005). and there is inevitably a selection bias depending on whether the sample is X-ray or optically selected.," The distribution of $\alpha_{OX}$ in normal QSOs is known to depend on luminosity \citep{vignali03,strateva05} and there is inevitably a selection bias depending on whether the sample is X-ray or optically selected."
+ “Phe most appropriate reference. distribution of ων in the literature is that presented in Stratevaetal.(2005) for a sample of 155 Sloan Digital Sky Survey QSOs with mecdium-deep ccoverage., The most appropriate reference distribution of $\alpha_{OX}$ in the literature is that presented in \citet{strateva05} for a sample of 155 Sloan Digital Sky Survey QSOs with medium-deep coverage.
+ However. since this distribution is based on an optically-selected sample. while our sample is X-ray selected. we have constructed. a reference distribution of aon For X-ray selected QSOs drawn from the International X-ray Optical Survey (RINOS: Alasonetαἱ. 2000).," However, since this distribution is based on an optically-selected sample, while our sample is X-ray selected, we have constructed a reference distribution of $\alpha_{OX}$ for X-ray selected QSOs drawn from the International X-ray Optical Survey \citep[RIXOS; ][]{mason00}."
+.. We restrict. the reference sample to those RINOS QSOs at z>1L. which corresponds to a similar luminosity range to our N-rav-absorbecd 09Ο sample. because RINOS has a similar N-rav. flux limit to the hhard spectrum: survey from which our targets. were," We restrict the reference sample to those RIXOS QSOs at $z>1$, which corresponds to a similar luminosity range to our X-ray-absorbed QSO sample, because RIXOS has a similar X-ray flux limit to the hard spectrum survey from which our targets were"
+Forbidden emission lines are important diagnostics of the mass loss [rom voung stellar objects (Cabrit et al.,Forbidden emission lines are important diagnostics of the mass loss from young stellar objects (Cabrit et al.
+ 1990: Hartigan. Edwards. Ghandour 1995).," 1990; Hartigan, Edwards, Ghandour 1995)."
+ In voung stars. the forbidden lines are strong in low-density mass oulllows. and are (vpically blue-shilted as the rec-shiftecl component is obscured by a circumstellar disk.," In young stars, the forbidden lines are strong in low-density mass outflows, and are typically blue-shifted as the red-shifted component is obscured by a circumstellar disk."
+ Spectroscopic survevs (e.g.llirthetal.1997:SolfandBohm1999:Pvo2003) have shown that forhbiclden lines in T Tauri stars sometimes have double peakecl velocity profiles. with both a high-velocity component (IVC) on the order of hundreds of kin/s and a low-velocitv component. (LVC) closer to (he rest velocity of the star.," Spectroscopic surveys \citep[e.g.][]{hir97, sol99, pyo03} have shown that forbidden lines in T Tauri stars sometimes have double peaked velocity profiles, with both a high-velocity component (HVC) on the order of hundreds of km/s and a low-velocity component (LVC) closer to the rest velocity of the star."
+ The survevs also demonstrate (hat forbidden lines have components (hat extend στον from the star ancl also that only emit [rom very near (he star., The surveys also demonstrate that forbidden lines have components that extend greatly from the star and also that only emit from very near the star.
+ In the stars with clear double peaked forbidden line profiles. (he IWC extends outwarel on aresecond scales. while the LWC does not. as pointed out by IHirthetal.(1997).," In the stars with clear double peaked forbidden line profiles, the HVC extends outward on arcsecond scales, while the LVC does not, as pointed out by \citet{hir97}."
+. The difficulty in determining the origins of the near-star forbidden line emission exists because it is offset not more than a lew tenths of an aresecond [rom the star 1997)., The difficulty in determining the origins of the near-star forbidden line emission exists because it is offset not more than a few tenths of an arcsecond from the star \citep{hir97}.
+. The source region may well lie within a lew stellar radii of the source. ancl would require a spatial resolution requirement of milliareseconds (o observe directly. not possible with present technology.," The source region may well lie within a few stellar radii of the source, and would require a spatial resolution requirement of milliarcseconds to observe directly, not possible with present technology."
+ Further. voung stars form in regions where dust. gas. and other sources ol light can obscure the star itself [rom view. so it is often difficult to follow embecdcec jets all Che wav {ο the star (Eislóffeletal.2000).," Further, young stars form in regions where dust, gas, and other sources of light can obscure the star itself from view, so it is often difficult to follow embedded jets all the way to the star \citep{eis00}."
+. Fortunately. most stars are in binary svstems. and we can use the kinematics of binary T Daun stars to indirectly probe the origin of (he forbidden lines that are not extended from the star.," Fortunately, most stars are in binary systems, and we can use the kinematics of binary T Tauri stars to indirectly probe the origin of the forbidden lines that are not extended from the star."
+ Highly eccentric binary stars vield large Doppler shifts that appear in any emission or absorption features (hat participate in the orbital motion of the two stars., Highly eccentric binary stars yield large Doppler shifts that appear in any emission or absorption features that participate in the orbital motion of the two stars.
+ If the voune binary stars have circumstellar disks. and the origin of the near forbidden lines is linked to (he disks. (hen we can expect to see orbital variation of the forbidden lime profiles in (ances with line profile changes in photospheric features of the stars themselves.," If the young binary stars have circumstellar disks, and the origin of the near forbidden lines is linked to the disks, then we can expect to see orbital variation of the forbidden line profiles in tandem with line profile changes in photospheric features of the stars themselves."
+ DQ Tau and UZ Tan-E are pre-imain sequence systems with signatures of active accretion disks. ie. classical T-Tauri stars.," DQ Tau and UZ Tau-E are pre-main sequence systems with signatures of active accretion disks, i.e. classical T-Tauri stars."
+ Both systems are spectroscopic binaries; DQ Tau has a period of 15.8 davs (Mathieuetal.L997) and UZ Tau-E a period of 19 davs 2002)., Both systems are spectroscopic binaries; DQ Tau has a period of 15.8 days \citep{mat97} and UZ Tau-E a period of 19 days \citep{pra02}.
+. Each svstem has a hiehlv eccentric orbit. wilh spectra containing strong forbidden lines.," Each system has a highly eccentric orbit, with spectra containing strong forbidden lines."
+ It is worth noting that the Forbidden line profiles of DQ Tau and UZ Tau-E lack distinct dual-velocity components., It is worth noting that the forbidden line profiles of DQ Tau and UZ Tau-E lack distinct dual-velocity components.
+ Indeed. DO Tau does not have much of an HINC at all.," Indeed, DQ Tau does not have much of an HVC at all."
+ UZ Tau-E did have a somewhat noticeable double peaked profile in Hartiganetal.(1995)... and also has a microjet of 0.2 arcseconds observed in [O I] 6300 slitless imaging (Πανασ. Edwards. Pierson 2004).," UZ Tau-E did have a somewhat noticeable double peaked profile in \citet{har95}, and also has a microjet of 0.2 arcseconds observed in [O I] 6300 slitless imaging (Hartigan, Edwards, Pierson 2004)."
+ The short periods and hieh eccentricities of these (wo svstenis made (hem ideal candidates for localizing the near-star emission with the method outlined above., The short periods and high eccentricities of these two systems made them ideal candidates for localizing the near-star emission with the method outlined above.
+We have obtained the first. sub-mm photometry of. CIUS 1915|105 at 350 Cllz (s50-pun). together with radio data on two separate epochs during a time when the source was mildly active.,"We have obtained the first sub-mm photometry of GRS 1915+105 at 350 GHz $\umu$ m), together with radio data on two separate epochs during a time when the source was mildly active."
+ At all frequencies a significant amount of variability was observed. with a roughly Hat spectral index. supporting the hypothesis that emission from radio to infrared wavelengths is dominated by svnchrotron emission.," At all frequencies a significant amount of variability was observed, with a roughly flat spectral index, supporting the hypothesis that emission from radio to infrared wavelengths is dominated by synchrotron emission."
+ Sealing with respect to both direct. VLBI imaging and radiomuminfrared time clelavs indicate that future sub-nimir observations have the potential to probe very. close to the base of the jet., Scaling with respect to both direct VLBI imaging and radio–mm–infrared time delays indicate that future sub-mm observations have the potential to probe very close to the base of the jet.
+ The authors are very grateful for the wizardry” of scheduling wv Graeme Watt at the JCALL. and for assistance by Richarel Ppestage. Rob νίκος. Wayne Lollanc and Vim Jenness in co-ordinating and reducing the data.," The authors are very grateful for the `wizardry' of scheduling by Graeme Watt at the JCMT, and for assistance by Richard Prestage, Rob Ivison, Wayne Holland and Tim Jenness in co-ordinating and reducing the data."
+ The Green Bank Interferometer is a [facility of the rational Science. Foundation operated by the NRAO in support of NASA Lligh I2nergy Astrophysics programmes., The Green Bank Interferometer is a facility of the National Science Foundation operated by the NRAO in support of NASA High Energy Astrophysics programmes.
+ tacdio astronomy at the Naval Research Laboratory is supported. by the Ollice of Naval Research., Radio astronomy at the Naval Research Laboratory is supported by the Office of Naval Research.
+ “Phe Ile Telescope is supported. by PPARC., The Ryle Telescope is supported by PPARC.
+ The Janes Clerk Alaxwell “Telescope is operated. by “Phe Joint Astrononi Centre on behalf of the Particle Physics anc Astronomy Research Council of the United Ixingdom. the Netherlands Organisation. for Scientific Research and the National Research Council of Canada.," The James Clerk Maxwell Telescope is operated by The Joint Astronomy Centre on behalf of the Particle Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organisation for Scientific Research and the National Research Council of Canada."
+Kamar (2000). if the observations are not frequent enough.,"Kumar (2000), if the observations are not frequent enough."
+" At the optical wavelength. the early light curve behaves tvpicallv as Fyx)77700n4UP yatheyr than F,x4.tt in the adiabatic case."," At the optical wavelength, the early light curve behaves typically as $F_{\nu}\propto t^{-(1+2\epsilon)/(4-\epsilon)}\sim t^{-0.45}$ rather than $F_{\nu}\propto t^{-1/4}$ in the adiabatic case."
+" When v,,, crosses the observing frequency. ie. ¢>/,,. Uae optical and X-ray light curves decays as £j,x4UP220/6000SpEEL as shown in cases D ancl A. Il should be noted (hat many X-ray alterglow light curves and a considerable fraction of optical alterglow light curves have temporal decaving indices sleeper than predicted (3p—2)/471.15 bv the standard adiabatic model (lor p22.2)."," When $\nu_{m}$ crosses the observing frequency, i.e. $t>t_{m}$, the optical and X-ray light curves decays as $F_{\nu}\propto
+t^{-(3p-2+2\epsilon)/(4-\epsilon)}\sim t^{-1.44}$, as shown in cases B and A. It should be noted that many X-ray afterglow light curves and a considerable fraction of optical afterglow light curves have temporal decaying indices steeper than predicted $(3p-2)/4\sim 1.15$ by the standard adiabatic model (for $p\approx2.2$ )."
+" The observed decaving indices in the X-ray afterglow light curves are (ay)=1.3340.38. while the median value of the observed X-ray spectral indices Jy. EF,xv. is ~1.05 (Berger. Kulkarni. Frail 2003: De Pasquale οἱ al."," The observed decaying indices in the X-ray afterglow light curves are $\langle\alpha_{X}\rangle=1.33\pm0.38$, while the median value of the observed X-ray spectral indices $\beta_{X}$, $F_{\nu}\propto \nu^{-\beta_{X}}$, is $\sim1.05$ (Berger, Kulkarni, Frail 2003; De Pasquale et al."
+ 2003)., 2003).
+ Assuming5 the X-ray frequency is above the cooling5 [requency and the spectral index is $y=p/2. the measured p is consistent with the standard value of the index of electron energy. distribution. p=2.2—2.3. predicted in the relativistic shock acceleration mechanism (see Achterberg οἱ al.," Assuming the X-ray frequency is above the cooling frequency and the spectral index is $\beta_{X}=p/2$, the measured $p$ is consistent with the standard value of the index of electron energy distribution, $p=2.2-2.3$, predicted in the relativistic shock acceleration mechanism (see Achterberg et al."
+ 2001 and reference therein)., 2001 and reference therein).
+ ILowever. the observed mean temporal decaving index (0) requires a relatively larger (p). 2.44. provided the shock is adiabatice.," However, the observed mean temporal decaying index $\langle\alpha_{X}\rangle$ requires a relatively larger $\langle
+p\rangle$, $\sim2.44$, provided the shock is adiabatic."
+ There are several caveats on the observations of the X-ray afterglows., There are several caveats on the observations of the X-ray afterglows.
+ First. the temporal behavior of the X-ray afterglow is hardly influenced by the equal arrival time surface effect. which will mix (he earlier light from high latitudes into (he present light.," First, the temporal behavior of the X-ray afterglow is hardly influenced by the equal arrival time surface effect, which will mix the earlier light from high latitudes into the present light."
+ This effect is important especially for high observing lrequencies. e.g. the optical and N-ray.," This effect is important especially for high observing frequencies, e.g. the optical and X-ray."
+ The profile of surface emissivitv of (he relativistic shock is ring-like in these high frequencies (Sari 1993)., The profile of surface emissivity of the relativistic shock is ring-like in these high frequencies (Sari 1998).
+" This will moderatelv slow down the decreasing rate of the afterglow after/,4 for the theoretical light curve A. However. the X-ray. afterglow is immune to this effect because its L,, is verv early aud (he observed decreasing index ay is based on (he observations (vpically several hours after the burst."," This will moderately slow down the decreasing rate of the afterglow after$t_{m}$ for the theoretical light curve A. However, the X-ray afterglow is immune to this effect because its $t_{m}$ is very early and the observed decreasing index $\alpha_{X}$ is based on the observations typically several hours after the burst."
+ Second. the measured oy is reliable since (he N-ray. absorption in the medium along the line of sight takes place at v<1 keV while the observing window is ~2—10 keV. Lastly. one should be cautious when interpreting the property of X-ray alterglow with the svnehrotvon radiation mechanism. because the X-ray alterglow may be contaminated by the svuchrotron-sel-Compton component.," Second, the measured $\beta_{X}$ is reliable since the X-ray absorption in the medium along the line of sight takes place at $\nu\lesssim 1$ keV while the observing window is $\sim2-10$ keV. Lastly, one should be cautious when interpreting the property of X-ray afterglow with the synchrotron radiation mechanism, because the X-ray afterglow may be contaminated by the synchrotron-self-Compton component."
+ However. there have so lar been onlv a few X-rav afterglows that were confirmed to have the IC components.," However, there have so far been only a few X-ray afterglows that were confirmed to have the IC components."
+ Therefore. the radiative corrections to the afterglow light curves must be taken into account based on the observations.," Therefore, the radiative corrections to the afterglow light curves must be taken into account based on the observations."
+ As can be seen in Figure L.. the light eurve at high frequeney. (t&vpe A and D) flattens when the afterglow enters the slow-cooling phase. />/44.," As can be seen in Figure \ref{fig:ISM:lightcurves}, , the light curve at high frequency (type A and B) flattens when the afterglow enters the slow-cooling phase, $t>t_{cm}$."
+" Al the transition time lon Lhe spectrum nearby the observing Irequeney changes [rom 4.<r4)""« vlov«Ww.«v.the⋅ expressions for− the⋅ fIux density ave⋅ the same. Le. F,⊏⇂↓⊐⊇↓⊇ =Ενveo""POThe"," At the transition time $t_{cm}$, the spectrum nearby the observing frequency changes from $\nu_{c}<\nu_{m}<\nu$ to $\nu_{m}<\nu_{c}<\nu$,while the expressions for the flux density are the same, i.e. $F_{\nu}=F_{\nu,\max}\nu_{m}^{(p-1)/2}\nu_{c}^{1/2}\nu^{-p/2}$."
+" while flattening of the light curve results from the Compton parameter Y in (he flux density. i.e. F,xY tsincey.x(lY)?zY7."," The flattening of the light curve results from the Compton parameter $Y$ in the flux density, i.e. $F_{\nu}\propto Y^{-1}$ , since $\nu_{c}\propto(1+Y)^{-2}\approx Y^{-2}$."
+ The Compton parameter Y in the slow-cooling phase decreases slowly. contrary (0 ils constaney in the earlier f[ast-cooling phase.," The Compton parameter $Y$ in the slow-cooling phase decreases slowly, contrary to its constancy in the earlier fast-cooling phase."
+ From equation (54)) lor Y in the slow-cooling phase aud adopting A= 0. the change of the temporal indexaround {ρω is Xa=(24e)(p—2)/[(4e)(4 p)]. which is shown inthe last segments of," From equation \ref{eqn:sc:spectrum}) ) for $Y$ in the slow-cooling phase and adopting $k=0$ , the change of the temporal indexaround $t_{cm}$ is $\Delta\alpha=(2+\epsilon)(p-2)/[(4-\epsilon)(4-p)]$ which is shown inthe last segments of"
+the /-band.,the $i$ -band.
+ Therefore bulge-dominated galaxies will have greater ROOFR50 ratios. while disk-dominated galaxies will have smaller values.," Therefore bulge-dominated galaxies will have greater $R90/R50$ ratios, while disk-dominated galaxies will have smaller values."
+" To determine the end-products of our PSG sample. we compare the peak-normalised distributions of concentration indices for our population of PSGs to the early- and late-type population of galaxies with log (M,) «10.5 M.."," To determine the end-products of our PSG sample, we compare the peak-normalised distributions of concentration indices for our population of PSGs to the early- and late-type population of galaxies with log $_{\star}$ ) $<10.5$ $_{\odot}$."
+ In agreement with 2.. we find that the division between early- and late-type galaxies is where the OO is approximately 2.6 times greater than the 50.," In agreement with \citet{strateva01}, we find that the division between early- and late-type galaxies is where the $R90$ is approximately 2.6 times greater than the $R50$."
+ The distribution of concentration indices for our PSG sample is very similar to that of the early-type galaxies., The distribution of concentration indices for our PSG sample is very similar to that of the early-type galaxies.
+ On the other hand. ? showed that early-type galaxies classified in such a manner may be contaminated by up to by edge-on spirals and that the SDSS parameter(ZracDczv: which gives the fraction of light fitted by a de Vaucouleurs profile) provides a better differentiation between the early- and late-type galaxies as typical early-type elliptical galaxies are traditionally characterised by a de Vaucouleurs profile and as such have an 70.5.," On the other hand, \citet{masters10} showed that early-type galaxies classified in such a manner may be contaminated by up to by edge-on spirals and that the SDSS parameter; which gives the fraction of light fitted by a de Vaucouleurs profile) provides a better differentiation between the early- and late-type galaxies as typical early-type elliptical galaxies are traditionally characterised by a de Vaucouleurs profile and as such have an $ > 0.5$."
+" As such. we show the distribution of for our ""green"" PSGs and low-mass early- and type sample in Figure 3.."," As such, we show the distribution of for our “green” PSGs and low-mass early- and late-type sample in Figure \ref{fracdevfig}."
+ The PSG distribution is shown by the gray-shaded histogram. while the striped histogram in the left-panel represents the distribution for the early-types and the distribution for late-types are represented by the striped histogram in the right panel.," The PSG distribution is shown by the gray-shaded histogram, while the striped histogram in the left-panel represents the distribution for the early-types and the distribution for late-types are represented by the striped histogram in the right panel."
+ A quantitative comparison using the KS test yields a KS probability of 0.68 that the distributions of PSGs and early-types are derived from the same parent sample. while the distributions of PSGs and late-types yield a KS probability of «0.001.," A quantitative comparison using the KS test yields a KS probability of 0.68 that the distributions of PSGs and early-types are derived from the same parent sample, while the distributions of PSGs and late-types yield a KS probability of $<$ 0.001."
+ Therefore. we tind that the structural stellar morphologies of the PSGs within the “green valley” appear to be more closely related to the morphologies of low-mass early-type galaxies even though star formation has only been truncated recently.," Therefore, we find that the structural stellar morphologies of the PSGs within the “green valley” appear to be more closely related to the morphologies of low-mass early-type galaxies even though star formation has only been truncated recently."
+ We derive the w+ colours of our sample using the (fromSDSSDR7:?) which are determined from the best fit of each galaxy profile to the linear combination of the exponential and the de Vaucouleurs profiles., We derive the $u-r$ colours of our sample using the \citep[from SDSS DR7; ][]{abazajian09} which are determined from the best fit of each galaxy profile to the linear combination of the exponential and the de Vaucouleurs profiles.
+ In addition. these magnitude measurements are corrected for dust attenuation using the models of ?..," In addition, these magnitude measurements are corrected for dust attenuation using the models of \citet{calzetti00}."
+ The “green valley” of our sample's colour distribution is defined to be within the nominal colour range of LS<nος2.3., The “green valley” of our sample's colour distribution is defined to be within the nominal colour range of $1.8<u-r<2.3$.
+ The stellar mass estimates for each galaxy are measured by fitting the five optical wavebands from SDSS to star formation history libraries generated from stellar models of ??..," The stellar mass estimates for each galaxy are measured by fitting the five optical wavebands from SDSS to star formation history libraries generated from stellar models of \citet{maraston98,maraston05}."
+ The uncertainties in our derived. stellar masses are dominated by the inherent uncertainties within the stellar models used., The uncertainties in our derived stellar masses are dominated by the inherent uncertainties within the stellar models used.
+ More details on the parameterization of the star formation histories and the fitting process can be found in ? and references within., More details on the parameterization of the star formation histories and the fitting process can be found in \citet{schawinski10a} and references within.
+ The top row of panels in Figure + shows the colour-stellar mass distribution of our PSG sample as red solid circles and the SDSS galaxy sample of a particular type is represented by the black solid contours., The top row of panels in Figure \ref{colourmass} shows the colour–stellar mass distribution of our PSG sample as red solid circles and the SDSS galaxy sample of a particular type is represented by the black solid contours.
+ Black dotted-lined contours represent the distribution of the entire sample regardless of galaxy-type and are shown for comparison purposes., Black dotted-lined contours represent the distribution of the entire sample regardless of galaxy-type and are shown for comparison purposes.
+ The bottom row of panels in Figure 4. shows the distribution of the fraction of PSGs to the number of galaxies (of a particular type) in a given colour-stellar mass bin as a solid orange-shaded contour map overlaid on the solid contours showing the distribution of the SDSS galaxies of a particular type., The bottom row of panels in Figure \ref{colourmass} shows the distribution of the fraction of PSGs to the number of galaxies (of a particular type) in a given colour–stellar mass bin as a solid orange-shaded contour map overlaid on the solid contours showing the distribution of the SDSS galaxies of a particular type.
+" As can be seen from the top row panels of Figure 4.. our PSG sample appears to be spread over a fairly large wr colour range even though a majority (945€) of the PSGs have stellar masses below the transition mass of log AZ,« 10.5 M. which separate the low-mass star-forming galaxies from the high-mass passively-evolving bulge-dominated galaxies (2).."," As can be seen from the top row panels of Figure \ref{colourmass}, our PSG sample appears to be spread over a fairly large $u-r$ colour range even though a majority ) of the PSGs have stellar masses below the transition mass of log $M_{\star} <$ 10.5 $_{\odot}$ which separate the low-mass star-forming galaxies from the high-mass passively-evolving bulge-dominated galaxies \citep{kauffmann03}."
+" In fact. we do not find any PSG with a log AL,> Τι M.."," In fact, we do not find any PSG with a log $M_{\star} >$ 11.5 $_{\odot}$."
+ Such a stellar mass limit in our PSG sample is more clearly illustrated by the number fractions of PSGs to the number of galaxies within a particular colour-stellar mass range (see the bottom panels of Figure +))., Such a stellar mass limit in our PSG sample is more clearly illustrated by the number fractions of PSGs to the number of galaxies within a particular colour–stellar mass range (see the bottom panels of Figure \ref{colourmass}) ).
+" Figure 5. shows the percentage of PSGs per 10°"" M. bins.", Figure \ref{fracerrsmassbin} shows the percentage of PSGs per $^{0.5}$ $_{\odot}$ bins.
+" The average percentage of PSGs in the log Ad, bins between 9.5 and 10.5 M. is ~S times greater than that between the log AZ, bins between 10.5 and [1.5 M. (with a 3o significance).", The average percentage of PSGs in the log $M_{\star}$ bins between 9.5 and 10.5 $_{\odot}$ is $\sim8$ times greater than that between the log $M_{\star}$ bins between 10.5 and 11.5 $_{\odot}$ (with a $\sigma$ significance).
+ The uncertainties in the stellar masses are dominated by the uncertainties in the stellar population models and can be up to 10°? M. in stellar masses (e.g...2).., The uncertainties in the stellar masses are dominated by the uncertainties in the stellar population models and can be up to $^{0.3}$ $_{\odot}$ in stellar masses \citep[e.g.\ ][]{conroy09}.
+ One possible reason for the lack of high mass PSGs is likely to be because our sample is restricted to a very local volume., One possible reason for the lack of high mass PSGs is likely to be because our sample is restricted to a very local volume.
+ As such. the probability of finding massive galaxies in such a local volume is much less than at higher redshifts.," As such, the probability of finding massive galaxies in such a local volume is much less than at higher redshifts."
+ Galaxies with log stellar masses greater than |1.5 account for only of the entire galaxy sample., Galaxies with log stellar masses greater than 11.5 account for only of the entire galaxy sample.
+ In addition. high-redshift surveys of E+A galaxies are not as sensitive to smaller Cand fainter) galaxies.," In addition, high-redshift surveys of E+A galaxies are not as sensitive to smaller (and fainter) galaxies."
+ As such. the," As such, the"
+BATSE instrument on theCORO satellite. in March. 1994.,"BATSE instrument on the satellite, in March 1994."
+ BATSE has detected six outbursts from the source between ALID. 48361 ancl 50231 (Scott. AL. private communication). though the one BATSE outburst plotted in Figure 4 was much more intense than the remaining five (Wilson ct al.," BATSE has detected six outbursts from the source between MJD 48361 and 50231 (Scott M., private communication), though the one BATSE outburst plotted in Figure 4 was much more intense than the remaining five (Wilson et al."
+ 1994: Scott. AL. private communication).," 1994; Scott M., private communication)."
+ Phe most recent. X-rav detection is a 0.1-Crab (2-12 keV) outburst detected by theRATE satellites All Sky Monitor (ASAT) between 1996 September 29 and October δ. consistent with the 186.5 clay period (Corbet Remillard 1996).," The most recent X-ray detection is a 0.1-Crab (2-12 keV) outburst detected by the satellite's All Sky Monitor (ASM) between 1996 September 29 and October 8, consistent with the 186.5 day period (Corbet Remillard 1996)."
+ The ASM data show a possible Hare at the previous preclictecl outburst epoch. but with far less significance.," The ASM data show a possible flare at the previous predicted outburst epoch, but with far less significance."
+ Dv design. the photometric svstem is most suitable or determining the stellar. parameters in early-type stars.," By design, the photometric system is most suitable for determining the stellar parameters in early-type stars."
+ Unfortunately. the Be star case is complicated by emission rom the circumstellar disc. so that the observed. indices come functions of both the stellar parameters and of the disc parameters.," Unfortunately, the Be star case is complicated by emission from the circumstellar disc, so that the observed indices become functions of both the stellar parameters and of the disc parameters."
+ Fabregat Reelero (1990)— determined ransfornis for the indices to correct for circumstellar clleets., Fabregat Reglero \shortcite{juan} determined transforms for the indices to correct for circumstellar effects.
+ “Phe transforms are based. upon the (La) xwameter. which has been shown to correlate closely with circumstellar continuum and. Balmer discontinuity emission," The transforms are based upon the $\alpha$ ) parameter, which has been shown to correlate closely with circumstellar continuum and Balmer discontinuity emission"
+Ossenkopf V. Stutzki J. Department of Astronomy. Stockholn A newborn protostar generates a fast and well jet. surrounded by a wider angle win,"Ossenkopf V. Stutzki J. Department of Astronomy, Stockholm University, Stockholm, Sweden } A newborn protostar generates a fast and well collimated jet, possibly surrounded by a wider angle wind."
+ the possiblymaterial drivesdensity (bow-)shocks travellt the surroundingejected high medium and traced vibrational lines at excitation temperatures of ar K. As a consequence. slower and cold (10-20 K) outflows are formed by swept-up material. usu by CO.," In turn, the ejected material drives (bow-)shocks travelling through the surrounding high-density medium and traced by $_2$ ro-vibrational lines at excitation temperatures of around 2000 K. As a consequence, slower and cold (10–20 K) molecular outflows are formed by swept-up material, usually traced by CO."
+ Shocks heat the gas and trigger several such as endothermic chemical reactions and ice αἱ sublimation or sputtering., Shocks heat the gas and trigger several processes such as endothermic chemical reactions and ice grain mantle sublimation or sputtering.
+ Several molecular speci significant enhancements in their abundances (se Dishoeck Blake ‘ardsas observed obse millimeter 1998))tow. sa number of byoutflow al. 1998:: Jorrgense," Several molecular species undergo significant enhancements in their abundances (see e.g., van Dishoeck Blake \cite{vanblake}) ), as observed by observations at millimeter wavelengths towards a number of outflows (Garay et al. \cite{garay};"
+nBachiller wavelengthsPérrez Gutiérrrez 1997. BP9 et al. 2007)).," Bachiller Pérrez Gutiérrrez 1997, BP97 hereafter; rgensen et al. \cite{jorge}) )."
+ The link between the gas ¢ at ~ IO K and the hot 2000 K shocked component 1 understanding how the protostellar wind transfers - and back to the ambient medium., The link between the gas components at $\sim$ 10 K and the hot 2000 K shocked component is crucial to understanding how the protostellar wind transfers momentum and energy back to the ambient medium.
+ In th the energy of the chemical o molecular understandingbow-shock is essential bevause it compositionrepres diagnostic tool for its The pow," In this context, the understanding of the chemical composition of a typical molecular bow-shock is essential bevause it represents a very powerful diagnostic tool for probing its physical conditions."
+erfulL1157 outflow. located at a probingdistance physicalestim between 250 (Looney et al. 2007)) and labo," The L1157 outflow, located at a distance estimated to be between 250 pc (Looney et al. \cite{looney}) )"
+ratory440 pc (V may be regardedpc as the ideal for ob: effects of shocks on the chemicallygas chemistry. outflows the of the so-called rich being(Bach 2001)).," and 440 pc (Viotti \cite{viotti}) ) may be regarded as the ideal laboratory for observing the effects of shocks on the gas chemistry, being the archetype of the so-called chemically rich outflows (Bachiller et al. \cite{bach01}) )."
+ The low-luminosity (4-11 Class ( IRAS20386+6751 drives a Le)powerful outflow associated with bow precessingshocks seen in et al. 1996)) a, The low-luminosity (4–11 $L_{\rm \sun}$ ) Class 0 protostar IRAS20386+6751 drives a precessing powerful molecular outflow associated with several bow shocks seen in CO (Gueth et al. \cite{gueth96}) )
+land in IR severalimages Etsló Neufeld et 2009))., and in IR $_2$ images (Davis Eislöfffel \cite{davis}; Neufeld et al. \cite{neufeld09}) ).
+BI In H»particular. the (Davis E bow-shock. called (Fig. 1)).," In particular, the brightest blue-shifted bow-shock, called B1 (Fig. \ref{maps}) ),"
+ has been interferometersbrightestextensive mm wit," has been extensively mapped with the PdB and VLA interferometers at mm- and cm-observations revealing a rich and clumpy structure, the clumps being located at the wall of the cavity with an arch-shape (Tafalla Bachiller \cite{tafabachi}; Gueth et al. \cite{fred98};"
+," Benedettini et al. \cite{milena07}, ,"
+ho, hereafter BVC07; Codella et al. \cite{code09}) ).
+bservationsthe PdBrevealing," L1157-B1 is well traced by molecules thought to be released by dust mantles such as $_2$ CO, $_3$ OH, and $_3$ as well as typical tracers of high-speed shocks such as SiO (e.g., Gusdorf et al. \cite{gus08b}) )."
+and a VLArich and c," Temperatures $\simeq$ 60–200 K (from $_3$, $_3$ CN, and SiO) as well as around 1000 K (from $_2$ ) have been derived (Tafalla Bachiller \cite{tafabachi}; Codella et al. \cite{code09};"
+l," Nisini et al. \cite{nisini07},"
+ump, in prep.).
+y structure.at located at the wall of th," However, a detailed study of the excitation conditions of the B1 structure has yet to be completed because of the limited range of excitation covered by the observations performed so far at cm- and mm-wavelengths."
+e with an bei, Observations of sub-mm lines with high excitation $\ge$ 50–100 K above the ground state) are thus required.
+ng(Tafalla 1995:: Gueth et al. cavity.," As part of the Key Program (Chemical Surveys of Star forming regions), L1157-B1 is currently being investigated with an unbiased spectral survey using the HIFI instrument (de Graauw et al. \cite{hifi}) )."
+1998:: Bei Y r 03 BachillerTAATPETIACAO py Pootgy," In this Letter, we report thefirst results based on HIFI observations in the 555–636 GHz spectral window, confirming the chemical richness and revealing different molecular components at different excitation conditions coexisting in the B1 bow structure."
+Tt las been suggested that black hole evaporation in braneworld eravity models cau lead to a deficiency of low-mass black holes in the steady state population.,It has been suggested that black hole evaporation in braneworld gravity models can lead to a deficiency of low-mass black holes in the steady state population.
+ A eap could be created in this context because the rate of evaporation iu braneworld eravity is very rapid aud ducreases with decreasing black hole mass (Postunov Chierepashcliuk 2003)., A gap could be created in this context because the rate of evaporation in braneworld gravity is very rapid and increases with decreasing black hole mass (Postnov Cherepashchuk 2003).
+ Towever. recent coustraimts ou the rate of evaporation obtained using the current population of black holes preclude this possibility (Johannsen. Psaltis. MeChliutock 2009).," However, recent constraints on the rate of evaporation obtained using the current population of black holes preclude this possibility (Johannsen, Psaltis, McClintock 2009)."
+ The paucity of black holes with masses less than 5AL. ave likely. therefore. to be related to the plivsics of superuova explosions that lead to the formation of black holes.," The paucity of black holes with masses less than $5~M_\odot$ are likely, therefore, to be related to the physics of supernova explosions that lead to the formation of black holes."
+ Numerical simulations of supernova explosions typically eeucrate a continuous distribution of black hole masses that decays as an exponential (Frver 1999: Frver Isalogera 2001)., Numerical simulations of supernova explosions typically generate a continuous distribution of black hole masses that decays as an exponential (Fryer 1999; Fryer Kalogera 2001).
+ The continuity of masses is primarily a consequence of the relatively eradual dependence of explosion energies ou the masses of the progenitors aud the fact that explosion cuereies are still comparable to (although smaller than) the binding euerev of the stellar cuvelopes., The continuity of masses is primarily a consequence of the relatively gradual dependence of explosion energies on the masses of the progenitors and the fact that explosion energies are still comparable to (although smaller than) the binding energy of the stellar envelopes.
+ Although creating a mass gap is difficult given the current understanding of the supernova energeties. it has beeu sugeested that it can be achieved under the ad hoc assumption that the explosion energv has a step-functiou dependence on progenitor mass and that it pluuges to zero for stars more massive than ~25M. (Fever Ixalogera 2001).," Although creating a mass gap is difficult given the current understanding of the supernova energetics, it has been suggested that it can be achieved under the ad hoc assumption that the explosion energy has a step-function dependence on progenitor mass and that it plunges to zero for stars more massive than $\sim
+25~M_\odot$ (Fryer Kalogera 2001)."
+ Such a bimodality iu the energies of explosions that form neutron stars versus black holes should become appareut iu the large sample of supernovae anticipated from the ongoing surveys that are sensitive to underbuninuous core-collapse sSuperiovae., Such a bimodality in the energies of explosions that form neutron stars versus black holes should become apparent in the large sample of supernovae anticipated from the ongoing surveys that are sensitive to underluminous core-collapse supernovae.
+ We thauk Chris Frver for stimulating discussions aud Rou Beuillaud for sharing his private catalogue of black hole candidates., We thank Chris Fryer for stimulating discussions and Ron Remillard for sharing his private catalogue of black hole candidates.
+ FO aud DP thank the ITC at the Iarviid-Siiithsouiui Ceuter for Astroplivsies for their hospitalitv., FÖ and DP thank the ITC at the Harvard-Smithsonian Center for Astrophysics for their hospitality.
+ FO acknowledges support frou NSF eraut. AST 07-08610 and Chaudra Theory. eraut. TMO-11003X. DP was supported by the NSF CAREER award NSF 0716519., FÖ acknowledges support from NSF grant AST 07-08640 and Chandra Theory grant TMO-11003X. DP was supported by the NSF CAREER award NSF 0746549.
+of cosmological simulations with either thermal or kinetic SN feedback.,of cosmological simulations with either thermal or kinetic SN feedback.
+ They showed that even when artificially preventing cooling of heated particles the results of the two schemes are considerably different., They showed that even when artificially preventing cooling of heated particles the results of the two schemes are considerably different.
+" Nevertheless, before tackling the over-cooling problem, one should first make sure that the input energy is accurately conserved."," Nevertheless, before tackling the over-cooling problem, one should first make sure that the input energy is accurately conserved."
+" In a recent study of the modelling of SN feedback in SPH simulations, Saitoh&Makino(2009) (hereafterSM09,seealsoMerlinetal.2010) noted that strong perturbations of the internal energy of gas particles lead to the violation of energy conservation when an individual particle time-step scheme is used (see Fig."," In a recent study of the modelling of SN feedback in SPH simulations, \cite{Saitoh2009} \citep[hereafter SM09, see
+also][]{Merlin2010} noted that strong perturbations of the internal energy of gas particles lead to the violation of energy conservation when an individual particle time-step scheme is used (see Fig."
+ 1 for an illustration)., \ref{fig:sedovindiv} for an illustration).
+ That can be summarised as follows., That can be summarised as follows.
+" If a source of energy alters the internal energy of a (active) gas particle, the particle will eventually decrease its time-step according to the Courant criterion."," If a source of energy alters the internal energy of a (active) gas particle, the particle will eventually decrease its time-step according to the Courant criterion."
+" However, if its neighbouring particles are they will not react immediately to the change of the thermal state of the region, and the integration accuracy may be strongly compromised."," However, if its neighbouring particles are they will not react immediately to the change of the thermal state of the region, and the integration accuracy may be strongly compromised."
+ The lack of a prompt response to the nearby energy injection can lead to effects like inter-particle crossing (due to particles missing the deceleration given by viscous forces) and non-conservation of energy., The lack of a prompt response to the nearby energy injection can lead to effects like inter-particle crossing (due to particles missing the deceleration given by viscous forces) and non-conservation of energy.
+" SM09 proved that, when using an appropriate time-step limiter, accurate energy conservation is achieved for strong internal energy perturbations."," SM09 proved that, when using an appropriate time-step limiter, accurate energy conservation is achieved for strong internal energy perturbations."
+ They proposed a scheme in which the ratio of the time-steps (longer over shorter) of neighbouring gas particles cannot be larger than a given factor fstep., They proposed a scheme in which the ratio of the time-steps (longer over shorter) of neighbouring gas particles cannot be larger than a given factor $f_{\rm step}$.
+" A value of fstep=4 is enough to ensure energy conservation to similar level of accuracy given by a global time-step integration scheme (Merlinetal.(2010) suggested the values fap=[4, 8])."," A value of $f_{\rm step}=4$ is enough to ensure energy conservation to similar level of accuracy given by a global time-step integration scheme \cite{Merlin2010} suggested the values $f_{\rm
+ step}=[4,8]$ )."
+" In this work and for the first time, we investigate under what conditions thermal and kinetic feedback schemes lead to qualitatively identical results."," In this work and for the first time, we investigate under what conditions thermal and kinetic feedback schemes lead to qualitatively identical results."
+" We first show that the feedback prescriptions are equivalent if the integration of the hydrodynamical equations is done over global, system time-steps."," We first show that the feedback prescriptions are equivalent if the integration of the hydrodynamical equations is done over global, system time-steps."
+" In order to do that, we compare test simulations of Sedov’s blast wave problem (Sedov1959;Landau&Lifshitz1959) to its analytic solution, finding very good agreement and energy conservation to within1%.."," In order to do that, we compare test simulations of Sedov's blast wave problem \citep{Sedov1959,Landau1959} to its analytic solution, finding very good agreement and energy conservation to within."
+" With the aim of generalising the results, we investigate whether one can achieve similar agreement and accuracy with the more computationally efficient individual time-step integration scheme."," With the aim of generalising the results, we investigate whether one can achieve similar agreement and accuracy with the more computationally efficient individual time-step integration scheme."
+ We implement a modified version of the SMO09 time-step limiter in the public version of (Springel2005) consistently taking into account the time-step synchronisation and the underlying leapfrog integrator., We implement a modified version of the SM09 time-step limiter in the public version of \citep{Springel2005} consistently taking into account the time-step synchronisation and the underlying leapfrog integrator.
+ We test the robustness and speed of the limiter with Sedov’s explosion problem injecting either thermal or kinetic energy at a random time., We test the robustness and speed of the limiter with Sedov's explosion problem injecting either thermal or kinetic energy at a random time.
+" We demonstrate that thermal and kinetic feedback methods are equivalent and energy is accurately conserved, provided that the system promptly responds to the energy perturbation."," We demonstrate that thermal and kinetic feedback methods are equivalent and energy is accurately conserved, provided that the system promptly responds to the energy perturbation."
+" With the application of the time-step scheme to galaxy formation and cosmological simulations in mind, we study a more realistic problem: the release of energy in the presence of density and pressure gradients."," With the application of the time-step scheme to galaxy formation and cosmological simulations in mind, we study a more realistic problem: the release of energy in the presence of density and pressure gradients."
+" We apply the scheme to an off-centre explosion in a self-gravitating gas halo, and show that the concordance of feedback methods and accurate energy conservation are still achieved even for the extreme case where all available energy (either thermal or kinetic) is injected into only one particle."," We apply the scheme to an off-centre explosion in a self-gravitating gas halo, and show that the concordance of feedback methods and accurate energy conservation are still achieved even for the extreme case where all available energy (either thermal or kinetic) is injected into only one particle."
+" In standard cosmological simulation practice the new integration approach presented in this paper could, in principle, require more simulation steps, but will mainly populate the intermediate time-bin levels with more particles."," In standard cosmological simulation practice the new integration approach presented in this paper could, in principle, require more simulation steps, but will mainly populate the intermediate time-bin levels with more particles."
+" T'herefore, we could expect a slight increase of the computational time for galaxy formation simulations, a price to pay for a better handling of feedback events."," Therefore, we could expect a slight increase of the computational time for galaxy formation simulations, a price to pay for a better handling of feedback events."
+" However, we still want to emphasise that, in idealised simulations like the ones presented in this work, the proposed scheme enables much faster integration than the standard, individual time-step scheme since it preserve a high energy conservation accuracy."," However, we still want to emphasise that, in idealised simulations like the ones presented in this work, the proposed scheme enables much faster integration than the standard, individual time-step scheme since it preserve a high energy conservation accuracy."
+ 'The paper is organised as follows., The paper is organised as follows.
+ We first show in Sec., We first show in Sec.
+ 2 that concordance of feedback methods is achieved with a global time-step scheme., \ref{sec:concordance} that concordance of feedback methods is achieved with a global time-step scheme.
+ We then apply our individual time-step algorithm to Sedov's blast wave test and to the centre explosion in a self-gravitating gas halo in Sec. 3.., We then apply our individual time-step algorithm to Sedov's blast wave test and to the off-centre explosion in a self-gravitating gas halo in Sec. \ref{sec:individual}.
+ We conclude in Sec. 4.., We conclude in Sec. \ref{sec:discuss}.
+" In all the simulations presented in this paper, we use the so-called system of units in which the units of velocity, length and mass are unity."," In all the simulations presented in this paper, we use the so-called system of units in which the units of velocity, length and mass are unity."
+concerned.,concerned.
+ Moreover. this iuerease seenis to be a power-law function (certainly uot a logarithmic one) of the uuuber of small particles.," Moreover, this increase seems to be a power-law function (certainly not a logarithmic one) of the number of small particles."
+ Iu assessing the sienificance of these numerical experiments. one ust notice that the variables involved have an exponential role.," In assessing the significance of these numerical experiments, one must notice that the variables involved have an exponential role."
+ The number of particles. I. appears iu the exponent of the Wilbert space dimension C aud πιοΊσα] conrputation of the Shannon eutropies 96) has a computational cost that scales as a power of A. depending ou implementation |10]..," The number of particles, $I$, appears in the exponent of the Hilbert space dimension ${\cal N}$, and numerical computation of the Shannon entropies $S(J)$ has a computational cost that scales as a power of ${\cal N}$ , depending on implementation \cite{etna}."
+ The rms for this paper have required thousauds of epu hours ou large clusters of processors., The runs for this paper have required thousands of cpu hours on large clusters of processors.
+ Far frou being a limitation of the technique. this computational complexity reveals the plivsical depth of the Alicki Faunes eutropies. as discussed in |10]..," Far from being a limitation of the technique, this computational complexity reveals the physical depth of the Alicki Fannes entropies, as discussed in \cite{etna}."
+ It also explains why these quantum eutropies have not (vet) received the attention they deserve. because they are so difficult to compute.," It also explains why these quantum entropies have not (yet) received the attention they deserve, because they are so difficult to compute."
+ Iu this paper we have introduced a multiparticle version of thecelebrated, In this paper we have introduced a multi–particle version of thecelebrated
+To understand galaxy evolution we need to understand the stellar properties of galaxies at high redshift.,To understand galaxy evolution we need to understand the stellar properties of galaxies at high redshift.
+ These galaxies include samples of mainly Lyman-break galaxies (LBGs: e.g. Steidel et al., These galaxies include samples of mainly Lyman-break galaxies (LBGs; e.g. Steidel et al.
+" 1996. 1999, Pettini et al."," 1996, 1999, Pettini et al."
+ 2001. Shapley et al.," 2001, Shapley et al."
+ 2003. Bunker et al.," 2003, Bunker et al."
+ 2004. Ouchi et al.," 2004, Ouchi et al."
+ 2004. Burgarella et al.," 2004, Burgarella et al."
+ 2007). but also other types of galaxies such as Ενα emitters (LAEs: e.g. Moller Warren 1993. Fynbo et al.," 2007), but also other types of galaxies such as $\alpha$ emitters (LAEs; e.g. ller Warren 1993, Fynbo et al."
+ 2002. Venemans et al.," 2002, Venemans et al."
+ 2007. Nilsson et al.," 2007, Nilsson et al."
+ 2009a) or very red and massive galaxies (DRGs. EROs; Franx et al.," 2009a) or very red and massive galaxies (DRGs, EROs; Franx et al."
+ 2003. Daddi et al.," 2003, Daddi et al."
+ 2004)., 2004).
+ One common method of studying the stellar populations of these galaxies is by fitting their spectral energy distributions (SEDs) with spectral templates. thereby determining Iuminosity-weighted properties such αν stellar ages. dust content. metallicities. and stellar masses.," One common method of studying the stellar populations of these galaxies is by fitting their spectral energy distributions (SEDs) with spectral templates, thereby determining luminosity-weighted properties such as stellar ages, dust content, metallicities, and stellar masses."
+ For Lyman-break galaxies. the last decade has seen a large number of publications determining the stellar properties using SED fitting.," For Lyman-break galaxies, the last decade has seen a large number of publications determining the stellar properties using SED fitting."
+ Two of the first publications. Shapley et al. (," Two of the first publications, Shapley et al. ("
+2001) and Papovich et al. (,2001) and Papovich et al. (
+2001) presented results for LBGs at :2. 23.,2001) presented results for LBGs at $z=2-3$ .
+ They found that LBGs at these redshifts tended to be young. with ages of a few hundred Myrs respectively. have masses around 1079 .. and significant extinction of the order of Ay=0.7.1.2 mag.," They found that LBGs at these redshifts tended to be young, with ages of a few hundred Myrs respectively, have masses around $10^{10}$ $_{\odot}$ and significant extinction of the order of $_V = 0.7 - 1.2$ mag."
+ The star formation rates (SFRs) in their samples were a few times 10M... yr.+., The star formation rates (SFRs) in their samples were a few times 10 $_\odot$ $^{-1}$.
+ SED fitting results at higher redshifts seem to indicate some evolution in these properties., SED fitting results at higher redshifts seem to indicate some evolution in these properties.
+ Verma et al. (, Verma et al. (
+2007) and Yabe et al. (,2007) and Yabe et al. (
+2009) fit LBGs at ;~5.5 and find that they are typically younger (a few times 10 Myrs). less massive (a few times 10? Μ.Ο). and less dusty (ντ0:3.0.7 mag.).,"2009) fit LBGs at $z \sim 5.5$ and find that they are typically younger (a few times 10 Myrs), less massive (a few times $10^9$ $_{\odot}$ ), and less dusty $A_V = 0.3 - 0.7$ mag.)."
+ They also find that the higher redshift LBGs have higher SFRs. of the order 50100 M.. yr.|. indicating that these are small galaxies in the process of building up their stellar mass.," They also find that the higher redshift LBGs have higher SFRs, of the order $50 - 100$ $_\odot$ $^{-1}$, indicating that these are small galaxies in the process of building up their stellar mass."
+ The evolution seen from smaller. younger galaxies at +~5 to larger and less vigorously star-forming galaxies at 2~2.5 does not seem to continue to even lower redshift.," The evolution seen from smaller, younger galaxies at $z\sim5$ to larger and less vigorously star-forming galaxies at $z \sim 2.5$ does not seem to continue to even lower redshift."
+ Overzier et al. (, Overzier et al. (
+"2009) presented a sample of local ""LBG analogues” based on UV bright galaxies in SDSS.",2009) presented a sample of local “LBG analogues” based on UV bright galaxies in SDSS.
+ These galaxies have masses of 6«10? M.. and SFRs of ~10M. yr+. ie. they are of similar mass and have similar SFRs to LBGs at :~2.5.," These galaxies have masses of $6 \times 10^9$ $_{\odot}$ and SFRs of $\sim 10$ $_\odot$ $^{-1}$, i.e. they are of similar mass and have similar SFRs to LBGs at $z \sim 2.5$."
+ Whereas SED fitting remains a popular method to determine stellar properties of LBGs. potential caveats with this method are the lack of a good wavelength spread in. or numbers of. spectral points and low signal-to-noise in the observations.," Whereas SED fitting remains a popular method to determine stellar properties of LBGs, potential caveats with this method are the lack of a good wavelength spread in, or numbers of, spectral points and low signal-to-noise in the observations."
+ In this paper we aim to avoid both these potential sources of uncertainty by observing a sample of pre-selected. lower redshift LBGs. at -~1 (Burgarella et al.," In this paper we aim to avoid both these potential sources of uncertainty by observing a sample of pre-selected, lower redshift LBGs, at $z\sim 1$ (Burgarella et al."
+ 2007)., 2007).
+ We also focus especially on LBGs with associated HST/ACS (45001 grism spectra. since the spectra will. at these redshifts. completely encompass the Balmer and 4000 breaks. thus allowing a very good sampling of these breaks.," We also focus especially on LBGs with associated HST/ACS G800L grism spectra, since the spectra will, at these redshifts, completely encompass the Balmer and 4000 breaks, thus allowing a very good sampling of these breaks."
+ These breaks. in turn. are very effective indicators of the age of the stellar population and dust content (Hamilton 1985. Balogh et al.," These breaks, in turn, are very effective indicators of the age of the stellar population and dust content (Hamilton 1985, Balogh et al."
+" 1999, Kriek et al."," 1999, Kriek et al."
+ 2006)., 2006).
+ In Kriek et al. (, In Kriek et al. (
+2006) it is shown that not only the redshift. but also the stellar properties are better constrained when using spectroscopic observations covering the Balmer and 4000 break of the galaxy. rather tha just broad-band photometric points.,"2006) it is shown that not only the redshift, but also the stellar properties are better constrained when using spectroscopic observations covering the Balmer and 4000 break of the galaxy, rather than just broad-band photometric points."
+ Kauffmann et al. (, Kauffmann et al. (
+2003) show that stellar properties can be determined solely on the basis of the 4000 break strength and the Balmer absorption-line index Hà4.,2003) show that stellar properties can be determined solely on the basis of the 4000 break strength and the Balmer absorption-line index $\delta_A$.
+ Slitless grism observations with the HST G800L grism are particularly well suited to study these breaks in high redshift galaxies: with low wavelength resolution. but high sensitivity and low background.," Slitless grism observations with the HST G800L grism are particularly well suited to study these breaks in high redshift galaxies; with low wavelength resolution, but high sensitivity and low background."
+Here. we present the results of fitting 15 LBGs with grism spectra. and show that the stellar properties are very well constrained by these observations.,"Here, we present the results of fitting 15 LBGs with grism spectra, and show that the stellar properties are very well constrained by these observations."
+ In Sect., In Sect.
+ 22 we describe the data used in the SED fitting. and in Sect.," \ref{sec:data} we describe the data used in the SED fitting, and in Sect."
+ ??the fitting method is described., \ref{sec:method} the fitting method is described.
+ The results are given in Sect., The results are given in Sect.
+ 22. and we end with a discussion in Sect. ??.., \ref{sec:results} and we end with a discussion in Sect. \ref{sec:disc}.
+" Throughout this paper. we assume a cosmology with 77τὸ km ! 1 O,,=0.3 and Q,= 0.7. Magnitudes are given in the AB system."," Throughout this paper, we assume a cosmology with $H_0=72$ km $^{-1}$ $^{-1}$ , $\Omega _{\rm m}=0.3$ and $\Omega _\Lambda=0.7$ Magnitudes are given in the AB system."
+picked because it is the wavelength least alfected by dust.,picked because it is the wavelength least affected by dust.
+ The extinelion correction is the major correction (ο the observed data., The extinction correction is the major correction to the observed data.
+ Extinction correction moves pixels horizontally to the right in figure 1 as can be seen bv the difference between the corrected and uncorrected points., Extinction correction moves pixels horizontally to the right in figure \ref{fig:hx} as can be seen by the difference between the corrected and uncorrected points.
+ The corrected h(x) must be a on the (rue (x) at low values of x for (vo reasons., The corrected h(x) must be a on the true h(x) at low values of x for two reasons.
+ First. pixels in a galaxy thal are below our detection limit due to extinction will not be included. and second. entire galaxies that [adl below our detection limit due to high extinction values will not be included.," First, pixels in a galaxy that are below our detection limit due to extinction will not be included and second, entire galaxies that fall below our detection limit due to high extinction values will not be included."
+ The break in the corrected data al log(x) = -3.5 is most probably due to the first effect. alühough lailure of the selection function at this level maa also contribute.," The break in the corrected data at log(x) = -3.5 is most probably due to the first effect, although failure of the selection function at this level may also contribute."
+ The faintest detected pixels lie at log(x) = -5.0. 30 times fainter than the break value.," The faintest detected pixels lie at log(x) = -5.0, 30 times fainter than the break value."
+ The difference between the faintest and the break value corresponds to an E(D-V) value of 0.4 which is toward the higher end of the exGinelions found lor the sample., The difference between the faintest and the break value corresponds to an E(B-V) value of 0.4 which is toward the higher end of the extinctions found for the sample.
+ Values of h(x) past log(x) = -3.5 should not be considered as accurate. however .the character of both the power law slope aud the exponential fall off are well established by the distribution at log(x) values greater (han -3.5.," Values of h(x) past log(x) = -3.5 should not be considered as accurate, however ,the character of both the power law slope and the exponential fall off are well established by the distribution at log(x) values greater than -3.5."
+ The observed distribution has two main components. a power law component at low x values and an exponential fall off at hieh x values.," The observed distribution has two main components, a power law component at low x values and an exponential fall off at high x values."
+ This Schechter like distribution (Press is verv common in astronomical phenomena., This Schechter like distribution \citep{sch74} is very common in astronomical phenomena.
+ If either the power law were extended to high values of x. or the exponential to low values of x. the total star formation rate would civerge.," If either the power law were extended to high values of x, or the exponential to low values of x, the total star formation rate would diverge."
+ Given (his familiar form of the distribution it is interesting to see if it arises naturally [rom known general empirical laws regarding mass distribution in galaxies and relations between gas density and star formation., Given this familiar form of the distribution it is interesting to see if it arises naturally from known general empirical laws regarding mass distribution in galaxies and relations between gas density and star formation.
+ Formulation in terms of these general relations can also give insight into whether and how the distribution Iunction might evolve with recshilt., Formulation in terms of these general relations can also give insight into whether and how the distribution function might evolve with redshift.
+ The derivation of the distribution shape utilizes (wo general empirical laws and an assumption., The derivation of the distribution shape utilizes two general empirical laws and an assumption.
+ The (wo empirical laws are the Schinidt law with a critical densitv relating star ormation intensity to gas surface clensitv and a Sehechter law for the distribution of galaxy nasses., The two empirical laws are the Schmidt law with a critical density relating star formation intensity to gas surface density and a Schechter law for the distribution of galaxy masses.
+ The assumption is that star formation occurs preclominantly in exponential disks., The assumption is that star formation occurs predominantly in exponential disks.
+ The following uses a climensionless mass variable y—min where m* is (he mass parameter in the Schechter mass distribution.," The following uses a dimensionless mass variable $y
+= m/m^*$ where $m^*$ is the mass parameter in the Schechter mass distribution."
+ ó* is given in m per comoving \Ipe?ni in eqn. 2.., $\phi^*$ is given in $m^*$ per comoving $^3$ in eqn. \ref{eqn:shc}.
+ In the ollowing mass alwavs refers (ο (he gas mass in the ealaxy since we are compuling star ormation., In the following mass always refers to the gas mass in the galaxy since we are computing star formation.
+ The formulation of the Schmidt Law comes from Kennicutt(1998) equ., The formulation of the Schmidt Law comes from \citet{ken98} eqn.
+ 7., 7.
+indicated by crosses.,indicated by crosses.
+ In this panel we find 10 white dwarfs. marked by open circles which would be sub-luminous if set at their kinematic distances.," In this panel we find 10 white dwarfs, marked by open circles which would be sub-luminous if set at their kinematic distances."
+ We conclude that these stars must be at farther distances from the Sun than predicted. therefore have higher tangential velocities. and must be rated as probable non-Hyades.," We conclude that these stars must be at farther distances from the Sun than predicted, therefore have higher tangential velocities, and must be rated as probable non-Hyades."
+ The best coincidence between the probable Hyades white dwarfs and the reference loci is seen in the left panel. one of the reasons being the better optical photometry available for the white dwarfs in the Johnson B.V system.," The best coincidence between the probable Hyades white dwarfs and the reference loci is seen in the left panel, one of the reasons being the better optical photometry available for the white dwarfs in the Johnson $B,V$ system."
+" In the near-infrared. NIR. panel (middle) the scatter in the /—K, colour is much larger than in B—V. because the fainter stars are at the detection limit of 2MASS. especially in the Κι band."," In the near-infrared, NIR, panel (middle) the scatter in the $ J-K_s$ colour is much larger than in $ B-V $, because the fainter stars are at the detection limit of 2MASS, especially in the $K_s$ band."
+ To a moderate extent. this holds for the 7’—/ colour. too.," To a moderate extent, this holds for the $ r'-J$ colour, too."
+" We find five stars in the middle panel (M, vs. J— Κι) that perfectly lie on the Hyades NIR main sequence.", We find five stars in the middle panel $M_J $ vs. $ J-K_s$ ) that perfectly lie on the Hyades NIR main sequence.
+ We marked these stars in Fig., We marked these stars in Fig.
+ 1. with their numbers from Table 1.., \ref{cmd} with their numbers from Table \ref{table:1}.
+ All five are included in our sample of 724 Hyades members in Paper I. Four of them (1. 4. 7 and 13 or HR 1358. EGGR 38. V471 Tau and WD 02174375. respectively) are known as binaries containing a WD and à MS component.," All five are included in our sample of 724 Hyades members in Paper I. Four of them (1, 4, 7 and 13 or HR 1358, EGGR 38, V471 Tau and WD 0217+375, respectively) are known as binaries containing a WD and a MS component."
+ The first three belong to the classical Hyades members. whereas WD 02174375 (13) has not be associated with the Hyades before.," The first three belong to the classical Hyades members, whereas WD 0217+375 (13) has not be associated with the Hyades before."
+ LP 649-0071 (12) is rated a white dwarf in Luyten’s White Dwarf Catalogues. though it has a NIR-colour typical of red dwarfs.," LP 649-0071 (12) is rated a white dwarf in Luyten's White Dwarf Catalogues, though it has a NIR-colour typical of red dwarfs."
+ This indicates a possible binary nature of this object., This indicates a possible binary nature of this object.
+ We discuss its properties in more detail in section 4.., We discuss its properties in more detail in section \ref{individual}.
+ In Table 1. we summarise the data for the white dwarfs discussed in this paper., In Table \ref{table:1} we summarise the data for the white dwarfs discussed in this paper.
+ Column 1 is a running number. column bo the name(s) of the star in the SIMBAD database. column 2 the spectral type taken from MeCook&Sion(1999)..," Column 1 is a running number, column 2 the name(s) of the star in the SIMBAD database, column 3 the spectral type taken from \citet{1999ApJS..121....1M}."
+ In column 4 we present the distance D of the star from the Sun às calculated from the convergent point method. whereas column 5 gives the distance r. from the cluster centre.," In column 4 we present the distance D of the star from the Sun as calculated from the convergent point method, whereas column 5 gives the distance $_c$ from the cluster centre."
+ Column 6 is the tangential velocity v_ perpendicular to the direction to the convergent point., Column 6 is the tangential velocity $_\perp$ perpendicular to the direction to the convergent point.
+ It is a measure of how well the motion of the star and cluster coincide., It is a measure of how well the motion of the star and cluster coincide.
+" Columns 7 to 11 give My. B—V. M,. J-K, and +’—J."," Columns 7 to 11 give $M_V$, $B-V$, $M_J$, $J-K_s$ and $r'-J$."
+ Column 12 describes whether or not the star is detected as an x-ray source (in VizieR). column 13 whether it is a known spectroscopic binary.," Column 12 describes whether or not the star is detected as an x-ray source (in VizieR), column 13 whether it is a known spectroscopic binary."
+ Finally. column 14 presents the sources of the B and V magnitudes.," Finally, column 14 presents the sources of the $B$ and $V$ magnitudes."
+ For code X.X we used the transformations from ugriz (SDSS) to B. V as given in Jesteretal.(2005)..," For code x,x we used the transformations from ugriz (SDSS) to $B$, $V$ as given in \citet{2005AJ....130..873J}."
+ An extended version of Table | is published only in machine-readable form via the CDS., An extended version of Table \ref{table:1} is published only in machine-readable form via the CDS.
+ It contains additional entries for each star including. e.g. precise positions. proper motions. apparent magnitudes to ease the preparation of follow-up observations. as well as the velocities derived from the convergent point method.," It contains additional entries for each star including, e.g. precise positions, proper motions, apparent magnitudes to ease the preparation of follow-up observations, as well as the velocities derived from the convergent point method."
+ The 37 stars in Table | are divided into four classes., The 37 stars in Table \ref{table:1} are divided into four classes.
+" The first ten stars of class ] are the ""classical"" Hyades.", The first ten stars of class 1 are the “classical” Hyades.
+ Stars 11 to 22 form class 2 of new probable Hyades co-movers., Stars 11 to 22 form class 2 of new probable Hyades co-movers.
+ The five stars of class 3 fulfil the kinematic and photometric criteria. but are more than 20 pe away from the centre in Z direction (perpendicular to the galactic plane).," The five stars of class 3 fulfil the kinematic and photometric criteria, but are more than 20 pc away from the centre in $Z$ direction (perpendicular to the galactic plane)."
+ Stars with these characteristics have been ruled out in Paper I. because all of them had discordant radial velocities (whenever a radial velocity measurement was available).," Stars with these characteristics have been ruled out in Paper I, because all of them had discordant radial velocities (whenever a radial velocity measurement was available)."
+ Finally. class 4+ consists of 10 stars that fulfil the kinematic criteria. but would be luminous in the CMD if set at their predicted distances.," Finally, class 4 consists of 10 stars that fulfil the kinematic criteria, but would be sub-luminous in the CMD if set at their predicted distances."
+ We also checked if measured radial velocities of the 37 candidates were available to compare them with the predicted ones from the convergent point method., We also checked if measured radial velocities of the 37 candidates were available to compare them with the predicted ones from the convergent point method.
+ Only for 12 of the 37 candidates we found radial velocities in the literature., Only for 12 of the 37 candidates we found radial velocities in the literature.
+ For six of them (the stars nos., For six of them (the stars nos.
+ 1. 4. 7. 8. ΤΙ. 21). the predicted and the measured radial velocities (from at least one source) agree well.," 1, 4, 7, 8, 11, 21), the predicted and the measured radial velocities (from at least one source) agree well."
+ EGGR 42 (10) was assumed to have Hyades radial velocity by Greenstein&Trimble(1967) and this was used to obtain its Einstein redshift (see also the remark on this star in the next section)., EGGR 42 (10) was assumed to have Hyades radial velocity by \citet{1967ApJ...149..283G} and this was used to obtain its Einstein redshift (see also the remark on this star in the next section).
+ WD 0816-387 (36) was already rejected as a Hyades member by the photometric criteria. so a disagreement between measured and predicted radial velocities is to be expected.," WD 0816+387 (36) was already rejected as a Hyades member by the photometric criteria, so a disagreement between measured and predicted radial velocities is to be expected."
+ On the other hand. the discordant radial velocities for the stars nos.," On the other hand, the discordant radial velocities for the stars nos."
+ 2. 19. 18. and 25 require a more detailed discussion.," 2, 13, 18, and 25 require a more detailed discussion."
+ Α reliable determination of space velocities of white dwarts is a challenging task., A reliable determination of space velocities of white dwarfs is a challenging task.
+ For isolated white dwarfs the apparent radial velocities must be corrected for gravitational redshift. which requires the knowledge of the mass-radius ratios. Le. quantities that cannot be observed directly.," For isolated white dwarfs the apparent radial velocities must be corrected for gravitational redshift, which requires the knowledge of the mass-radius ratios, i.e. quantities that cannot be observed directly."
+ For stars nos., For stars nos.
+ 2. 8. and 18. Paulietal.(2006) determined radial velocities from high-resolution spectra. whereas spectroscopic distances and gravitational. redshifts were computed from the fundamental parameters derived by Koesteretal.(2001)..," 2, 8, and 18, \citet{2006A&A...447..173P} determined radial velocities from high-resolution spectra, whereas spectroscopic distances and gravitational redshifts were computed from the fundamental parameters derived by \citet{2001A&A...378..556K}."
+ The relatively high radial velocity for EGGR 26 (2) by Paulietal.(2006) would reject this star as a Hyades member. though in numerous studies its membership is found to be confirmed (e.g.Weidemannetal.1992:DeGennaro2009)..," The relatively high radial velocity for EGGR 26 (2) by \citet{2006A&A...447..173P} would reject this star as a Hyades member, though in numerous studies its membership is found to be confirmed \citep[e.g.,][]
+{1992AJ....104.1876W,2009ApJ...696...12D}."
+ The discrepancy for EGGR 26 may probably be explained by underestimated uncertainties introduced when deriving the redshift corrections., The discrepancy for EGGR 26 may probably be explained by underestimated uncertainties introduced when deriving the redshift corrections.
+ The same reason for discrepancy. may possibly hold for LP 653-0026 (18). too: recently. Koesteretal.(2009) published an updated version of their catalogue of the fundamental parameters of white dwarfs where two different sets of parameters were considered to be equally probable for this star.," The same reason for discrepancy may possibly hold for LP 653-0026 (18), too: recently, \citet{2009A&A...505..441K}
+ published an updated version of their catalogue of the fundamental parameters of white dwarfs where two different sets of parameters were considered to be equally probable for this star."
+ A re-calculation of the radial velocities seems to be reasonable for white dwarfs from Paulietal.(2006).., A re-calculation of the radial velocities seems to be reasonable for white dwarfs from \citet{2006A&A...447..173P}.
+ For four of our candidates (the stars nos., For four of our candidates (the stars nos.
+ 7. 11. 13. 25). radial velocities were obtained by Silvestrietal.(2002) from line-of-sight velocities of M dwarfs in common proper notion pairs (cpm). each consisting of an M-dwarf and a white dwarf.," 7, 11, 13, 25), radial velocities were obtained by \citet{2002AJ....124.1118S}
+ from line-of-sight velocities of M dwarfs in common proper motion pairs (cpm), each consisting of an M-dwarf and a white dwarf."
+ The authors assume that typical separations between the components are about 1000 AU. such that orbital motion can be neglected.," The authors assume that typical separations between the components are about 1000 AU, such that orbital motion can be neglected."
+ For star no., For star no.
+ 13. where measured and predicted radial velocities differ by 1.6c. the separation is2”. corresponding to about 50 AU. given a distance of 25 pe.," 13, where measured and predicted radial velocities differ by $\sigma$, the separation is, corresponding to about 50 AU, given a distance of 25 pc."
+ Here the orbital notion cannot be neglected. and even a small correction of a few kms! could make the difference between measured and predicted radial velocities insignificant.," Here the orbital motion cannot be neglected, and even a small correction of a few $\rm{km\,s}^{-1}$ could make the difference between measured and predicted radial velocities insignificant."
+ On the other hand. if the separation is large in à cpm pair. the argument of a common radial velocity becomes weaker because of an increasing probability of an unphysical optical pair.," On the other hand, if the separation is large in a cpm pair, the argument of a common radial velocity becomes weaker because of an increasing probability of an unphysical optical pair."
+ This could be the case of star no., This could be the case of star no.
+ 25. WD 06414438. where the separation between the white dwarf and its MS companion reaches 143..," 25, WD 0641+438, where the separation between the white dwarf and its MS companion reaches ."
+ The PPMXL lists more or less compatible motions for these, The PPMXL lists more or less compatible motions for these
+We apply the same proof (word by word) of Lemma {92 and Corollary L23.. with the origin 0 shifted to fQ and the time step Af replaced by foty.,"We apply the same proof (word by word) of Lemma \ref{lem1} and Corollary \ref{corr1}, with the origin $0$ shifted to $t_{1}$ and the time step $\D t$ replaced by $t_{2}-t_1$."
+ LI Now we are ready to prove Proposition L.1.., $\hfill{\Box}$ Now we are ready to prove Proposition \ref{bef_proof}.
+ From the definition (1.17)) of ο. we know that The case where eq=(Af)0 is obvious.," From the definition \ref{ebar}) ) of $\bar{e}$, we know that The case where $e_{1}=\bar{e}(\D t)=0$ is obvious."
+ Four cases could be0)., Four cases could be.
+ Tn this case we have cy=64. c=ei aud ¢y=0.," In this case we have $e_{0}=e^{+}_{0}$, $e_{1}=e^{+}_{1}$ and $e^{-}_{0}=0$."
+ Therefore (1.2)) is auiuniediate consequence of (1.15)).0)., Therefore \ref{sh7mra}) ) is animmediate consequence of \ref{eqn_cor1}).
+ Similar to Case 0).In this case e4=6!(At)., Similar to Case .In this case $e_{1}=e^{+}(\D t)$.
+" Let the time £fj be defined as follows: Using inequality CLI8)) with £;=f ,.f2=Ataud Af—Ff|,<At. we get: where we have used the fact that ο(ty)=e(f4)= 0. aud heuce (1.2)) O}.."," Let the time $t_{-+}$ be defined as follows: Using inequality \ref{79bas}) ) with $t_{1}=t_{-+}$, $t_{2}=\D t$ and $\underline{\D} t=\D t-t_{-+}\leq
+\D t$, we get: where we have used the fact that $e^{+}(t_{1})=e^{-}(t_{1})=0$ , and hence \ref{sh7mra}) )."
+ Similar to Case 3., Similar to Case 3.
+ LI This section is entirely devoted to the proof of Theorcim 1.5.., $\hfill{\Box}$ This section is entirely devoted to the proof of Theorem \ref{theo2}.
+ Let T>0 and let Af>0 be such that where » to be chosen large enough (the choice of Af will be given later)., Let $T>0$ and let $\D t>0$ be such that where $n$ to be chosen large enough (the choice of $\D t$ will be given later).
+ Tn order to estimate |u*alle~<cu.T)- we add aud subtract e. the continuous piecewise linear function passing through the points (CEA. Oy k=0n. Iu other words where eand Aj ave defined iu CL1)).," In order to estimate $\|u^{\eps}-u^{0}\|_{L^{\infty}(0,T)}$, we add and subtract$v$ , the continuous piecewise linear function passing through the points $(k\D t,
+v^{k})$ , $k=0\cdots n$ In other words where $v^{k}$and $\l_{k}$ are defined in \ref{dis_sch}) )."
+ We start by stating the following corollary that eeneralizes Proposition LA., We start by stating the following corollary that generalizes Proposition \ref{bef_proof}. .
+" EF,~LO?teres. Ly,~I0eres. (M~0.01—10, ", $E_{\gamma} \sim 10^{51} \mathrm{ergs} $ $E_\mathrm{tot} \sim 10^{52} \mathrm{ergs}$ $ \dot{M} \sim 0.01-10 M_{\odot}$ 
+measured in the infrared.,measured in the infrared.
+ For NGC 2808 we got time series for 11 stars and a single measurement for two further objects., For NGC 2808 we got time series for 11 stars and a single measurement for two further objects.
+ Only three variables could be observed in NGC 362., Only three variables could be observed in NGC 362.
+ Photometry on the K-band images was performed in the same way as for the visual data., Photometry on the $K$ -band images was performed in the same way as for the visual data.
+ The zero point of the magnitude scale was set using 4 to 8 stars on the frame that had 2MASS magnitudes., The zero point of the magnitude scale was set using 4 to 8 stars on the frame that had 2MASS magnitudes.
+ Observed magnitudes were converted to the 2MASS system by using the relation given in Carpenter (2001))., Observed magnitudes were converted to the 2MASS system by using the relation given in Carpenter \cite{Carpenter01}) ).
+ All individual infrared measurements for each variable (detected from visual data) were combined to get an average K magnitude., All individual infrared measurements for each variable (detected from visual data) were combined to get an average $K$ magnitude.
+ Using 2 nonvariable 2MASS sources we estimate an accuracy of the mean K brightness of mmag., Using 2 nonvariable 2MASS sources we estimate an accuracy of the mean $K$ brightness of mag.
+ In both clusters a number of long period variables were found., In both clusters a number of long period variables were found.
+ We identified the variables in the 2MASS point source catalogue where possible., We identified the variables in the 2MASS point source catalogue where possible.
+ With the exception of two stars (1 in NGC 362 and 1 in NGC 2808) close to the cluster centre we could in this way attribute coordinates and 2MASS photometry data to our targets., With the exception of two stars (1 in NGC 362 and 1 in NGC 2808) close to the cluster centre we could in this way attribute coordinates and 2MASS photometry data to our targets.
+ Four known and seven new long period variables were found in NGC 362., Four known and seven new long period variables were found in NGC 362.
+" For a large fraction of the variables in this cluster we had to rely on the ANDICAM data only for determining the period, because for the MACHO observations the central part of the cluster unfortunately fell onto damaged areas of the detector for a significant fraction of the time series."," For a large fraction of the variables in this cluster we had to rely on the ANDICAM data only for determining the period, because for the MACHO observations the central part of the cluster unfortunately fell onto damaged areas of the detector for a significant fraction of the time series."
+" The stars with ANDICAM data only are identified in refn362vars,, where coordinates, V, Ic, 2MASS magnitudes and periods are listed."," The stars with ANDICAM data only are identified in \\ref{n362vars}, where coordinates, $V$, $I_C$, 2MASS magnitudes and periods are listed."
+ As in Lebzelter Wood (2005)) we used the variable star identifiers given by Clement et al. (2001)), As in Lebzelter Wood \cite{LW05}) ) we used the variable star identifiers given by Clement et al. \cite{Clement97}) )
+ where available., where available.
+" All other variables were named ""LW"" with some number."," All other variables were named ""LW"" with some number."
+" Most of these variables have been detected for the first time, see below for a few exceptions."," Most of these variables have been detected for the first time, see below for a few exceptions."
+ The location of the variables in a CMD of the cluster can be seen in reffhd362.., The location of the variables in a CMD of the cluster can be seen in \\ref{fhd362}.
+ The phased light curves are presented in refn362lc.., The phased light curves are presented in \\ref{n362lc}.
+ It is quite obvious that most stars show some kind of irregular behaviour on top of the adopted periodic light change., It is quite obvious that most stars show some kind of irregular behaviour on top of the adopted periodic light change.
+ For some targets a reasonable fit was only possible by adding a long secondary period like in the case of LW8., For some targets a reasonable fit was only possible by adding a long secondary period like in the case of LW8.
+" However, an accurate period for these long time variations could not be determined due to the limited length of our time series."," However, an accurate period for these long time variations could not be determined due to the limited length of our time series."
+" In addition to the variables listed here we noticed three other red stars that showed some kind of light variability, but the variations did not follow any (semi-)periodicity."," In addition to the variables listed here we noticed three other red stars that showed some kind of light variability, but the variations did not follow any (semi-)periodicity."
+" These stars are listed at the bottom of Table 2,, but were not investigated further."," These stars are listed at the bottom of Table \ref{n362vars}, but were not investigated further."
+ Periods exist in the literature for two stars of our NGC 362 sample., Periods exist in the literature for two stars of our NGC 362 sample.
+ As can be seen we confirm the period of dd for V2 originally found by Sawyer (1931))., As can be seen we confirm the period of d for V2 originally found by Sawyer \cite{SH31}) ).
+" It is not clear, if the star is changing period from time to time or if the dd period of Eggen (1972)) resulted from a short time irregularity in the lightcurve?."," It is not clear, if the star is changing period from time to time or if the d period of Eggen \cite{Eggen72}) ) resulted from a short time irregularity in the light."
+. The literature period for V16 is also found in our data., The literature period for V16 is also found in our data.
+ Our star LW1 is identical with the star V36 of Szekely et al. (2007))., Our star LW1 is identical with the star V36 of Szekely et al. \cite{Szekely07}) ).
+ We confirm the suspicion that the star is a long period variable., We confirm the suspicion that the star is a long period variable.
+" A second star from the list of Szekely et al.,"," A second star from the list of Szekely et al.,"
+" V56, is located very close to our variable LW6."," V56, is located very close to our variable LW6."
+" However, these authors classified it as an eclipsing binary."," However, these authors classified it as an eclipsing binary."
+" While we cannot completely reject this hypothesis, we think that a classification as LPV seems to be more probable."," While we cannot completely reject this hypothesis, we think that a classification as LPV seems to be more probable."
+ Infrared data were added to Table 2 from the 2MASS catalogue., Infrared data were added to Table \ref{n362vars} from the 2MASS catalogue.
+" As mentioned above our ANDICAM near-IR photometry provided multiple K-magnitudes for only three sources from our list (LW7, LW8 and LW9)."," As mentioned above our ANDICAM near-IR photometry provided multiple $K$ -magnitudes for only three sources from our list (LW7, LW8 and LW9)."
+" LW7 has a mean K=8.72x0.05 mag, for LW8 we find K=9.11+0.05 mag."," LW7 has a mean $K=8.72\pm0.05$ mag, for LW8 we find $K=9.11\pm0.05$ mag."
+" In both cases the scatter is above the expected photometric accuracy, but no reasonable light curve could be extracted."," In both cases the scatter is above the expected photometric accuracy, but no reasonable light curve could be extracted."
+ For LW9 we derived an average K magnitude of mmag., For LW9 we derived an average $K$ magnitude of mag.
+" The star shows light variations with a K amplitude of up to mmag, which are rather well synchronized with the observed change in the visual."," The star shows light variations with a $K$ amplitude of up to mag, which are rather well synchronized with the observed change in the visual."
+ Our K-band photometry is in good agreement with the 2MASS data., Our $K$ -band photometry is in good agreement with the 2MASS data.
+" Two investigations dedicated to cluster membership of stars in NGC 362 are found in the literature, namely the proper motion study by Tucholke (1992)) and the radial velocity study of Fischer et al. (1993))."," Two investigations dedicated to cluster membership of stars in NGC 362 are found in the literature, namely the proper motion study by Tucholke \cite{Tucholke92}) ) and the radial velocity study of Fischer et al. \cite{Fischer93}) )."
+ Proper motion data are also found in the NOMAD database (Zacharias et al. 2005))., Proper motion data are also found in the NOMAD database (Zacharias et al. \cite{Zacharias05}) ).
+" Five stars from our sample have been measured by Tucholke: V2, LW8 and LW11 are listed as proper motion members of the cluster."," Five stars from our sample have been measured by Tucholke: V2, LW8 and LW11 are listed as proper motion members of the cluster."
+" LW1 and LW9 have a low probability for membership, but the proper motion data for theses two stars come with a large error bar."," LW1 and LW9 have a low probability for membership, but the proper motion data for theses two stars come with a large error bar."
+" LW8, LWIO and LW11 through LW13"," LW8, LW10 and LW11 through LW13"
+associated with the Sevfert nucleus.,associated with the Seyfert nucleus.
+ Removing NGC 1068. we can form a new ACN composite spectrum from the remaining three ACN galaxies.," Removing NGC 1068, we can form a new AGN composite spectrum from the remaining three AGN galaxies."
+ Although: HON was not detected iu auv of the individual galaxies. it is reaclily detected in their composite spectrum.," Although HCN was not detected in any of the individual galaxies, it is readily detected in their composite spectrum."
+ In this new composite AGN spectrum. we fiud ratios of ICO! /TICN = 3.12:0.9 and 00 = 15430.12.," In this new composite AGN spectrum, we find ratios of $^+$ /HCN = $\pm$ 0.9 and $^{13}$ CO = $\pm$ 0.12."
+ Not surprising. these ratios are even imore extreme relative to the starburst sample than the AGN composite includiug NGC 10683.," Not surprising, these ratios are even more extreme relative to the starburst sample than the AGN composite including NGC 1068."
+ The line ratios discussed iu the previous section depend on both the phivsieal (density. column density. and temperature) aud chemical (auolecular abundance) xoperties of the gas.," The line ratios discussed in the previous section depend on both the physical (density, column density, and temperature) and chemical (molecular abundance) properties of the gas."
+ The information available iu our spectra is inadequate to constrain all the needed gas xoperties., The information available in our spectra is inadequate to constrain all the needed gas properties.
+ We have consequently modeled the Lue ratios with a very simplistic approach in which we assume hat the cussion arises from a collection of ideutical uolecular cloud cores which are characterized by a single cluperature and density. with equal abundances of the various molecular species;," We have consequently modeled the line ratios with a very simplistic approach in which we assume that the emission arises from a collection of identical molecular cloud cores which are characterized by a single temperature and density, with equal abundances of the various molecular species."
+ Molecular abundances were aken from the well studied Galactic CAIC cores MIT and Ceph A sumunarized by Bereinetal.(1997)., Molecular abundances were taken from the well studied Galactic GMC cores M17 and Ceph A summarized by \citet{ber97}.
+".. We ‘ther assume abundance ratios of CO/IL, = 1x10! aud PCO//CO = 24107: the resulting molecular abundances relative to Πο are given in Table |.", We further assume abundance ratios of $_2$ = $\times10^{-4}$ and /CO = $\times10^{-2}$; the resulting molecular abundances relative to $_2$ are given in Table 4.
+ Au E/A ratio of 1 for CILOII and an ortho/para ratio of 3 for Ca4llo were also assunued., An E/A ratio of 1 for $_3$ OH and an ortho/para ratio of 3 for $_3$ $_2$ were also assumed.
+ The assmmed molecular abundances are very similar to those derived by Martinet2006) for NCC 253 based ou their 2 nuu wavelength spectral survey (see comparison in Table 1)., The assumed molecular abundances are very similar to those derived by \citet{mar06} for NGC 253 based on their 2 mm wavelength spectral survey (see comparison in Table 4).
+ The cussion frou an eusenmible of GALC cores was modeled with the RADEN code (VanderTakctal. using the Einstein A-cocficients aud collision rate coefficients from the Leiden Atomic aud Molecular Database (Schóieretal.2005)., The emission from an ensemble of GMC cores was modeled with the RADEX code \citep{van07} using the Einstein -coefficients and collision rate coefficients from the Leiden Atomic and Molecular Database \citep{sch05}.
+. To compute the line intensities. besides the molecular abundauces. the gas density. ον gas temperature. Z. aud the Ilo cobuuu deusitv per unit line width. N(IL))/Ac. need to be specified.," To compute the line intensities, besides the molecular abundances, the gas density, $n$, gas temperature, $T$, and the $_2$ column density per unit line width, $_2)$ $\Delta$$v$, need to be specified."
+ Our standard CALC core model assumes T — 35 K. 2) = 14105 «n7. aud (Πο) Λος 14107 cin? (n 4) +.," Our standard GMC core model assumes $T$ = 35 K, $n$ = $\times$ $^5$ $^{-3}$, and $_2)$ $\Delta$$v$ = $\times$ $^{22}$ $^{-2}$ (km $^{-1}$ $^{-1}$."
+" Line widths for CAMICS in the Galactic disk are around 3 km  (Sucllctal.1981)...leading to a IL, column deusity of 3«1072 7."," Line widths for GMCs in the Galactic disk are around 3 km $^{-1}$ \citep{sne84},leading to a $_2$ column density of $\times$ $^{22}$ $^{-2}$."
+ Molecular. clouds near the ceuter of our Galaxy. such as Ser D2. have line widths around 20 kgu ? (Ciuminsetal.LO8G).. aud for this line width the modeled Ty column deusity is < 1078 cu>.," Molecular clouds near the center of our Galaxy, such as Sgr B2, have line widths around 20 km $^{-1}$ \citep{cum86}, and for this line width the modeled $_2$ column density is $\times$ $^{23}$ $^{-2}$."
+ The line ratios computed from our staudard model are compared with the observatious of NGC 253. the ealaxy with the richest spectrum. in Figure 3.," The line ratios computed from our standard model are compared with the observations of NGC 253, the galaxy with the richest spectrum, in Figure 3."
+ Since IICN usually has the strougest enuüssionu of any of the ‘high dipole moment® molecules. all iuteusity ratios are computed relative to ICN.," Since HCN usually has the strongest emission of any of the 'high dipole moment' molecules, all intensity ratios are computed relative to HCN."
+" Multiple lines were observed for several molecular species and for the line ratio analysis the 87.9252 Cz liue for IENC'O, the 96.7lL (11 line for CIT;OII aud the 99.2999 Cz line for SO were used iu foriung the intensity ratios."," Multiple lines were observed for several molecular species and for the line ratio analysis the 87.9252 GHz line for HNCO, the 96.7414 GHz line for $_3$ OH and the 99.2999 GHz line for SO were used in forming the intensity ratios."
+ For IIC4N. the two limes at 81.5815 and. 100.0764 CGIIz. which were in the least line coufused reeious of the spectrum. were sununed together.," For $_3$ N, the two lines at 81.8815 and 100.0764 GHz, which were in the least line confused regions of the spectrum, were summed together."
+ The line ratios for our standard core model agree well with those observed for NGC 253., The line ratios for our standard core model agree well with those observed for NGC 253.
+ The largest discrepancies between the observed aud modeled line ratios is less than a factor of two., The largest discrepancies between the observed and modeled line ratios is less than a factor of two.
+ The observed ratios in this ealaxy are well fit using this siuple model., The observed ratios in this galaxy are well fit using this simple model.
+ Qur model eau also be used to address the seusitivitv of the various line ratios to density aud temperature., Our model can also be used to address the sensitivity of the various line ratios to density and temperature.
+ The line ratios were computed from the standard model for densities of 1« 10! and 1109 cm2 and the resulting line ratios shown iu Figure 3., The line ratios were computed from the standard model for densities of $\times$ $^4$ and $\times$ $^6$ $^{-3}$ and the resulting line ratios shown in Figure 3.
+ For our assumed abunudanuces. a deusitv of «10? cin? is the better fit to the observed line ratios i NGC 253.," For our assumed abundances, a density of $\times$ $^5$ $^{-3}$ is the better fit to the observed line ratios in NGC 253."
+ Figure 3 illustrates that several of the line ratios. due to differences in the critical density of the lines formune the ratio. are liehly scusitive to density.," Figure 3 illustrates that several of the line ratios, due to differences in the critical density of the lines forming the ratio, are highly sensitive to density."
+ As expected. one of the values most affected is the rratio: however. in addition. the Bue ratios CILOIT/IICN and ΠΝΟΟ/ΠΟΝ are. also. density. seusitive.," As expected, one of the values most affected is the ratio; however, in addition, the line ratios $_3$ OH/HCN and HNCO/HCN are also density sensitive."
+ The temperature schsitivity of these ratios can also be explored. aud iu Figure 1 the observed line ratios in NGC 2523 are compared to those derived from our standard model with deusitv 110 cm° aud with varving temperatures.," The temperature sensitivity of these ratios can also be explored, and in Figure 4 the observed line ratios in NGC 253 are compared to those derived from our standard model with density $\times$ $^5$ $^{-3}$ and with varying temperatures."
+ Temperature has a mich smaller effect on the line ratios than docs deusity. largely because he molecular trausitious we observed all arise from simular energy levels above the erouud rotational state.," Temperature has a much smaller effect on the line ratios than does density, largely because the molecular transitions we observed all arise from similar energy levels above the ground rotational state."
+ NGC 253 las been widely studied aud there are a iniuber of estimates of the deusity of the molecular eas., NGC 253 has been widely studied and there are a number of estimates of the density of the molecular gas.
+ Based on fitting multiple lines of CS. IICN aud ICO! ed Martínetal.(2005) and IWamdsenetal.(2007) to estimate a deusity of approximately 2410? cu? veher than an earlicr estimate based on Πως ο emissiou (Thittemeisteretal.1997).," Based on fitting multiple lines of CS, HCN and $^+$ led \citet{mar05} and \citet{knu07} to estimate a density of approximately $\times$ $^5$ $^{-3}$, higher than an earlier estimate based on $_2$ CO emission \citep{hut97}."
+. Recently. Bavetetal.(2009) Bt the CS cinission from NCC 253 with a two-compoucut uodel cousisting of a cool. lower deusity component and a warm. higher density component.," Recently, \citet{bay09} fit the CS emission from NGC 253 with a two-component model consisting of a cool, lower density component and a warm, higher density component."
+ For the two velocity features In NGC 253. they found densities for he cooler component of order «105 and for the warmer compoucut of order 24109 >.," For the two velocity features in NGC 253, they found densities for the cooler component of order $\times$ $^4$ and for the warmer component of order $\times$ $^6$ $^{-3}$."
+ The density in our standard model agrees well with the single deusitv fits of Martinetal.(2005). and Ruudseuctal.(2007) and is consisteut with the average of the coupoucuts ft bv Bavetetal.(2009)., The density in our standard model agrees well with the single density fits of \citet{mar05} and \citet{knu07} and is consistent with the average of the components fit by \citet{bay09}.
+.. Based on this density. the abundances we asstuned in our standard model. iucludiug hose aolecules (INCO aud CIT;OIT) that may be xoduced in shocks. are consistent with the observatious in NGC 253.," Based on this density, the abundances we assumed in our standard model, including those molecules (HNCO and $_3$ OH) that may be produced in shocks, are consistent with the observations in NGC 253."
+" A better determination of the physical xoperties of the molecular gas is needed to refine the abundance estimates,", A better determination of the physical properties of the molecular gas is needed to refine the abundance estimates.
+ Our model can be used to examine the anti-correlation οπου the ICO! /IHCN aud ΠΟΝ ratios shown iu Figure 2., Our model can be used to examine the anti-correlation between the $^+$ /HCN and /HCN ratios shown in Figure 2.
+ Overlaid on the data in Figure 2 is shown a ine connecting the results for five models iu which only he deusity is varied. from 14104 to «109 cum7.," Overlaid on the data in Figure 2 is shown a line connecting the results for five models in which only the density is varied, from $\times$ $^4$ to $\times$ $^6$ $^{-3}$."
+ The conrparison between data aud models suggests that wach of the variation observed in the MCO//TICN ratio cau vc explained if the average density of the molecular gas varies απο these galaxies., The comparison between data and models suggests that much of the variation observed in the /HCN ratio can be explained if the average density of the molecular gas varies among these galaxies.
+ For deusitics below <1 cm ος it is difficult to produce detectable ICN or ccluission without having unrealistically laree cobuun deusities of these molecular species., For densities below $\times$ $^4$ $^{-3}$ it is difficult to produce detectable HCN or emission without having unrealistically large column densities of these molecular species.
+ Although the ICO! /IICN ratio is not strouglv affected by: density. soie of the auti-correlation found between this ratio aud the PCO/AIICN ratio can be. explained by variations in the average deusity.," Although the $^+$ /HCN ratio is not strongly affected by density, some of the anti-correlation found between this ratio and the /HCN ratio can be explained by variations in the average density."
+ Juneauetal.(2009) came to, \citet{jun09} came to
+The simplest wav to combine (he emission and absorpton spectra is (o ignore the blending ancl compute (he excitation temperature for each velocity channel naively assuming all the gas al a given. velocity. is at the same temperature.,The simplest way to combine the emission and absorption spectra is to ignore the blending and compute the excitation temperature for each velocity channel naively assuming all the gas at a given velocity is at the same temperature.
+ This gives the simple formula: where 75 is the brightness temperature of (he emission in velocity channel e. and το) is the optical depth of the absorption in the corresponding velocity channel.," This gives the simple formula: where $T_B$ is the brightness temperature of the emission in velocity channel $v$ , and $\tau(v)$ is the optical depth of the absorption in the corresponding velocity channel."
+ This gives the harmonic mean of the temperatures of (he various neutral atomic regions which contribute to that velocity channel. ie. where / is an index for the clifferent regions on the line of sight contributing. each with column density .V; and excitation temperature 7;.," This gives the harmonic mean of the temperatures of the various neutral atomic regions which contribute to that velocity channel, i.e. where $i$ is an index for the different regions on the line of sight contributing, each with column density $N_i$ and excitation temperature $T_i$."
+" The distribution of 7,4, vs. number of channels lor (he nine sources with 7, 730Ilx is given on figure 11.", The distribution of $T_{sp}$ vs. number of channels for the nine sources with $T_c >$ 30K is given on figure 11.
+ We must assume (hat (he temperatures resulting from equation 7 shown on figure 11 are strongly. biased to higher values than the true cool temperature due to blending with the warm gas., We must assume that the temperatures resulting from equation 7 shown on figure 11 are strongly biased to higher values than the true cool temperature due to blending with the warm gas.
+ ]nuproved estimates of the cool phase temperature are based on (he shapes of the features in (he emission and absorption spectra., Improved estimates of the cool phase temperature are based on the shapes of the features in the emission and absorption spectra.
+ A simple approach is to fit the absorption spectrum wilh a iumnber of Gaussian components., A simple approach is to fit the absorption spectrum with a number of Gaussian components.
+" Low latitude A21-cnm emission spectra do not generally look like the sum of distinct. Gaussian components. but absorption spectra clo, so (his is a reasonable approach."," Low latitude $\lambda$ 21-cm emission spectra do not generally look like the sum of distinct Gaussian components, but absorption spectra do, so this is a reasonable approach."
+ When the Gaussian components blend together the best fil parameters become ambiguous. ancl for a complicated. low latitude spectrum even the number of Gaussian components needed is not clear.," When the Gaussian components blend together the best fit parameters become ambiguous, and for a complicated, low latitude spectrum even the number of Gaussian components needed is not clear."
+ Thus there is no unique solution to the Gaussian fitting problem., Thus there is no unique solution to the Gaussian fitting problem.
+ The fitted parameters are obtained using an interactive program based on the Levenbere- method of non-linear chi-squared minization (Press et al..," The fitted parameters are obtained using an interactive program based on the Levenberg-Marquardt method of non-linear chi-squared minization (Press et al.,"
+ 1992)., 1992).
+ The numbered boxes above (he spectra on figures 2-8 indicate the velocity ranges dominated bx the different line components (given on Table 2. columns6 and 7).," The numbered boxes above the spectra on figures 2-8 indicate the velocity ranges dominated by the different line components (given on Table 2, columns6 and 7)."
+ Weaker components may be hidden, Weaker components may be hidden
+"When 5=1/2. When n,= 0. When 52—-1/2. Finally. when 5=-1.","When $n_B=1/2$, When $n_B=0$ , When $n_B=-1/2$ Finally, when $n_B=-1$,"
+and detected. hard X-ray. emission from the radio galaxy as well as the cluster emission.,and detected hard X-ray emission from the radio galaxy as well as the cluster emission.
+ We present the highly disturbed X-rav morphology of the cluster obtained from SCA and the ROSAT LURE and the A-rayv spectral properties. as evidence for a cluster merger and then discuss a possible interpretation on the cluster evolution., We present the highly disturbed X-ray morphology of the cluster obtained from ASCA and the ROSAT HRI and the X-ray spectral properties as evidence for a cluster merger and then discuss a possible interpretation on the cluster evolution.
+ The presence of an obscured active nucleus in 36353. is also revealed. through the X-ray spectrum. obtained with SCA., The presence of an obscured active nucleus in 3C353 is also revealed through the X-ray spectrum obtained with ASCA.
+ No redshift measurements have been available for the member galaxies apart from 3€353 (2= 0.030)., No redshift measurements have been available for the member galaxies apart from 3C353 $z=0.030$ ).
+ In this paper. we present a new optical CCD image of the central region of the cluster and the first redshift measurements for the member galaxies.," In this paper, we present a new optical CCD image of the central region of the cluster and the first redshift measurements for the member galaxies."
+ Zw l1Yls.]0105 and 830353 were observed. with ASCA and the ROSAT LIBI., Zw 1718.1–0108 and 3C353 were observed with ASCA and the ROSAT HRI.
+ A summary of the ASCA and ROSAT observations are given in Table 1., A summary of the ASCA and ROSAT observations are given in Table 1.
+ In addition to a point-like X-ray source at the position of the radio galaxy. a bright. extended N-rav. source is. detected. at. the South-East: of 30353 with the AACA Gas Imaging Spectrometer (CLS: C62 and G3).," In addition to a point-like X-ray source at the position of the radio galaxy, a bright, extended X-ray source is detected at the South-East of 3C353 with the ASCA Gas Imaging Spectrometer (GIS; G2 and G3)."
+ A likely source of the extended: X-ray. emission is the Zwicky cluster 1718.1.0108 of which 30353 is a member., A likely source of the extended X-ray emission is the Zwicky cluster 1718.1–0108 of which 3C353 is a member.
+ Since the primary target of the observations was 30353. the Solid state Imaging Spectrometer (S18: SO and S1) of ASCA was operating using the standard 1 CCD chip with a H1.11 arcmin field of view which covers the radio galaxy and only a small fraction (north part) of the cluster emission of Zw lY18.10105.," Since the primary target of the observations was 3C353, the Solid state Imaging Spectrometer (SIS; S0 and S1) of ASCA was operating using the standard 1 CCD chip with a $11\times 11$ arcmin field of view which covers the radio galaxy and only a small fraction (north part) of the cluster emission of Zw 1718.1–0108."
+ For the ASCA cata. standard calibration (used for the Revision 2 processing) and data reduction techniques were employed. using F'TOOLS (version 4.1). provided. by. the ASCA Guest. Observer Facility at Goddard Space Flight Center.," For the ASCA data, standard calibration (used for the Revision 2 processing) and data reduction techniques were employed, using FTOOLS (version 4.1) provided by the ASCA Guest Observer Facility at Goddard Space Flight Center."
+ The Cs data were mainly used to investigate the extended X-ray emission. from the cluster., The GIS data were mainly used to investigate the extended X-ray emission from the cluster.
+ Phe SIS data are used for investigating 3353. since it provides a better spatial resolution than the CIS. which helps to separate 30353 from the cilfuse cluster emission.," The SIS data are used for investigating 3C353, since it provides a better spatial resolution than the GIS, which helps to separate 3C353 from the diffuse cluster emission."
+ The ROSAT LHigh Resolution Imager (Pfellermann οἱ al 1987) eave a 0.12.4 keV image at the spatial resolution of about 5 aresec., The ROSAT High Resolution Imager (Pfeffermann et al 1987) gave a 0.1–2.4 keV image at the spatial resolution of about 5 arcsec.
+ With 17.2 ks of net exposure. a total of 5271 counts were detected.," With 17.2 ks of net exposure, a total of 5271 counts were detected."
+ Phe raw ΕΠ image has been smoothed with the adaptive kernel method. ASMOOTII (Ebeling. White Rangarajan 1999). using a gaussian kernel ancl a characteristic smoothing threshold. above the local background. of 2¢.," The raw HRI image has been smoothed with the adaptive kernel method, ASMOOTH (Ebeling, White Rangarajan 1999), using a gaussian kernel and a characteristic smoothing threshold, above the local background, of $2\sigma$."
+ Three X-ray clumps were detected., Three X-ray clumps were detected.
+ llowever. we note that the short LRRD exposure is not appropriate to study the cdilluse. low surface brightness rav emission of the cluster.," However, we note that the short HRI exposure is not appropriate to study the diffuse, low surface brightness X-ray emission of the cluster."
+ The full-band. CIS image superposed by the ROSAT HU contours is shown in Fig., The full-band GIS image superposed by the ROSAT HRI contours is shown in Fig.
+ 1., 1.
+" Phe angular scale is 2:0.05 Alpe tat the redshift of 0.03 (44,0 = 50 ty).", The angular scale is $\simeq 0.05$ Mpc $^{-1}$ at the redshift of 0.03 0 = 50 ).
+ The X-rav source is extended. roughly in the North-South direction: 30353 is also detected. at the NW. edge. of the cluster emission., The X-ray source is extended roughly in the North-South direction; 3C353 is also detected at the NW edge of the cluster emission.
+ Asvnumetric. slightlv twisted. X-ray morphology is evident in the GLS image.," Asymmetric, slightly twisted X-ray morphology is evident in the GIS image."
+ Phree X-ray peaks (N. €. and S from north to south) are resolved in the IRI image.," Three X-ray peaks (N, C, and S from north to south) are resolved in the HRI image."
+ Faint plumes extending to the NorthWest. (NWp) ancl West (Wop) are seen in the GIS image., Faint plumes extending to the NorthWest (NWp) and West (Wp) are seen in the GIS image.
+ The NW. plume is also seen at a faint level in the HEB image., The NW plume is also seen at a faint level in the HRI image.
+ The orientation of these X-ray components are given in Table 2., The orientation of these X-ray components are given in Table 2.
+ The two GIS images in the 0.72 keV and 5.510 keV bands are shown in Fig., The two GIS images in the 0.7–2 keV and 5.5–10 keV bands are shown in Fig.
+ 2., 2.
+ 3€353 is the brightest. source in the hard-band. image but faint in the κορα image. indicating a very hard spectrum. probably due to heavy obscuration in the active nucleus (see Section 4.2).," 3C353 is the brightest source in the hard-band image but faint in the soft-band image, indicating a very hard spectrum probably due to heavy obscuration in the active nucleus (see Section 4.2)."
+ We note that the hard. energy band. is similar to that (510. keV), We note that the hard energy band is similar to that (5–10 keV)
+m]] observations.,] observations.
+" In Figure 1 we show sample LOS spectra of u]], 2CO, CO, and CO. We observed the J=1—0 transitions of ""CO, CO, and C!80 toward the observed π]] LOSs."," In Figure \ref{fig:example_los} we show sample LOS spectra of ], $^{12}$ CO, $^{13}$ CO, and $^{18}$ O. We observed the $J = 1 \to 0$ transitions of $^{12}$ CO, $^{13}$ CO, and $^{18}$ O toward the observed ] LOSs."
+ These observations are part of a survey of all positions towards the inner Galaxy between {= —175.5?and |= 56.8°conducted at the ATNF Mopra Telescope., These observations are part of a survey of all positions towards the inner Galaxy between $l=-175.5$ and $l=56.8$ conducted at the ATNF Mopra Telescope.
+" The angular resolution of these observations is 33"".", The angular resolution of these observations is .
+". Typical system temperatures were 600, 300, and KK for 1200, 1900, and 0:50, respectively."," Typical system temperatures were 600, 300, and K for $^{12}$ CO, $^{13}$ CO, and $^{18}$ O, respectively."
+ To convert from antenna to main-beam temperature scale we use a main-beam efficiency of 0.42 (?).., To convert from antenna to main–beam temperature scale we use a main-beam efficiency of 0.42 \citep{Ladd05}.
+ All lines were observed simultaneously using the MOPS spectrometerin zoom mode., All lines were observed simultaneously using the MOPS spectrometerin zoom mode.
+" The spectra were smoothed in velocity to kkm s! for ""CO and CO and to ss! for C!*O. Typical rms noise is KK for ?CO and KK for both CO and 0/50. We checked pointing accuracy every 60 minutes using the closest and brightest SiO maser.", The spectra were smoothed in velocity to km $^{-1}$ for $^{12}$ CO and $^{13}$ CO and to $^{-1}$ for $^{18}$ O. Typical rms noise is K for $^{12}$ CO and K for both $^{13}$ CO and $^{18}$ O. We checked pointing accuracy every 60 minutes using the closest and brightest SiO maser.
+ We identify a total of 146 π]] velocity components in the observed LOSs., We identify a total of 146 ] velocity components in the observed LOSs.
+ From this data set we identify components that are associated with high-column density molecular gas by looking for CO counterparts., From this data set we identify components that are associated with high–column density molecular gas by looking for $^{13}$ CO counterparts.
+ We identified most of the high-'CO column density u]] components by fitting Gaussian functions defined by fitting the corresponding CO spectra., We identified most of the $^{13}$ CO column density ] components by fitting Gaussian functions defined by fitting the corresponding $^{13}$ CO spectra.