ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2" A clear trend is nevertheless observed, indicating higher extinction along the line of sight to image B, as suggested by Turnshek et al. (1997))."," A clear trend is nevertheless observed, indicating higher extinction along the line of sight to image B, as suggested by Turnshek et al. \cite{tur97}) )."3" The flux ratio is reasonably well modeled using where AAg=Ag(AÀ)-Ag(B) is the difference of extinction between the A and B line of sights measured in the B filter and &(A) the extinction curve tabulated in Pei (1992)), redshifted to the lens redshift z;«1.0 (Kneib et al. 1998))."," The flux ratio is reasonably well modeled using where $\Delta A_B = A_B ({\rm4A}) - A_B ({\rm B})$ is the difference of extinction between the A and B line of sights measured in the $B$ filter and $\xi(\lambda)$ the extinction curve tabulated in Pei \cite{pei92}) ), redshifted to the lens redshift $z_l \simeq 1.0$ (Kneib et al. \cite{kne98}) )."5" Since no obvious ffeature typical of the Milky Way extinction is observed at 2.3 um""! (Fig. 4)),"," Since no obvious feature typical of the Milky Way extinction is observed at 2.3 $\mu$ $^{-1}$ (Fig. \ref{fig:ratio}) ),"6 we adopt a SMC-like extinction curve., we adopt a SMC-like extinction curve.7 A reasonably good fit is obtained with FAo/Fgo = 1.05+0.02 and AAg = —0.09+0.02 (Fig. 4))., A reasonably good fit is obtained with $F_{\rm A0} / F_{\rm B0}$ = $\pm$ 0.02 and $\Delta A_B$ = $-$ $\pm$ 0.02 (Fig. \ref{fig:ratio}) ).8 With z;~1.88 (Goicoecha and Shalyapin 2010)) a similar fit is obtained with AAg = —0.055+0.015., With $z_l \simeq 1.88$ (Goicoecha and Shalyapin \cite{goi10}) ) a similar fit is obtained with $\Delta A_B$ = $-$ $\pm$ 0.015.9" For z;« 1.0 (1.88), the lines of sight to images A and B sample regions of the lens galaxy located at respectively 4.9 (5.2) kpc and 6.0 (6.4) kpc from the galaxy center, adopting the astrometry of Chantry and Magain (2007)) and assuming a flat cosmology with ©,, = 0.27 and Ho = 70 km s! Mpc!."," For $z_l \simeq$ 1.0 (1.88), the lines of sight to images A and B sample regions of the lens galaxy located at respectively 4.9 (5.2) kpc and 6.0 (6.4) kpc from the galaxy center, adopting the astrometry of Chantry and Magain \cite{cha07}) ) and assuming a flat cosmology with $\Omega_m$ = 0.27 and $_{\rm 0}$ = 70 km $^{-1}$ $^{-1}$ ."10 The flux ratio FAo/Fpgo = 1.19+0.10 measured by MacLeod et al. (2009))," The flux ratio $F_{\rm11A0} / F_{\rm B0}$ = $\pm$ 0.10 measured by MacLeod et al. \cite{mac09}) )"12" at 11 um in the mid-infrared, i.e. at wavelengths were both extinction and microlensing are expected to be negligible, is compatible with the extinction corrected flux ratio FAo/Fpgo we derive, although marginally higher."," at 11 $\mu$ m in the mid-infrared, i.e. at wavelengths were both extinction and microlensing are expected to be negligible, is compatible with the extinction corrected flux ratio $F_{\rm A0} / F_{\rm B0}$ we derive, although marginally higher."13" Time variations in the absorption line profiles of BAL quasars are not uncommon (Barlow et al. 1989,1992,,"," Time variations in the absorption line profiles of BAL quasars are not uncommon (Barlow et al. \cite{bar89,bar92},"14 Gibson et al. 2008)). InH1413+, Gibson et al. \cite{gib08}) ).15"117, Turnshek et al. (1988))"," In, Turnshek et al. \cite{tur88}) )"16 reported a deepening of the BAL between 1981 and 1985., reported a deepening of the BAL between 1981 and 1985.17" On the contrary, between 1989 and 2005 (Fig. 2))"," On the contrary, between 1989 and 2005 (Fig. \ref{fig:specvis}) ),"18" variations are observed as a gradual decrease of the depth of the BAL high-velocity part, the deepest component of the profile being essentially unaffected."," variations are observed as a gradual decrease of the depth of the BAL high-velocity part, the deepest component of the profile being essentially unaffected."19" Variations appear more complex in than iniv, affecting a larger part of the absorption profile (see also Fig. 6))."," Variations appear more complex in than in, affecting a larger part of the absorption profile (see also Fig. \ref{fig:sepavis}) )."20" The strongest change occurs between 1989 and 1993 and corresponds to an increase of the luminosity (Remy et al. 1996,,"," The strongest change occurs between 1989 and 1993 and corresponds to an increase of the luminosity (Remy et al. \cite{rem96},"21 Ostensen et al. 1997))., stensen et al. \cite{ost97}) ).22" Moreover, stronger absorption is accompanied by stronger emission, which is an indication that resonance line scattering can play an important role in the emission line formation."," Moreover, stronger absorption is accompanied by stronger emission, which is an indication that resonance line scattering can play an important role in the emission line formation."23" In the line of the AB spectrum, the high-velocity absorption appears ~15% larger in 1989 than in 2005 while the emission is ~25% more intense."," In the line of the AB spectrum, the high-velocity absorption appears $\sim$ larger in 1989 than in 2005 while the emission is $\sim$ more intense."24" In the framework of resonance scattering where each absorbed photon is re-emitted, this may suggest that the high-velocity outflow has more scattering material perpendicular to the line of sight than absorbing material along the line of sight."," In the framework of resonance scattering where each absorbed photon is re-emitted, this may suggest that the high-velocity outflow has more scattering material perpendicular to the line of sight than absorbing material along the line of sight."25 This also requires a large covering factor., This also requires a large covering factor.26 We follow a method similar to that used in Sluse et al. (2007))., We follow a method similar to that used in Sluse et al. \cite{slu07}) ).27" Assuming thatthe observed spectra F; of the different images are made of a superposition of a spectrum Fy which is only macrolensed and of a spectrum Fy, which is both macro- and", Assuming thatthe observed spectra $F_i$ of the different images are made of a superposition of a spectrum $F_M$ which is only macrolensed and of a spectrum $F_{M\mu}$ which is both macro- and28simulation models based on the computer will more closely fit the realistic physical situation.,simulation models based on the computer will more closely fit the realistic physical situation.29" Additionally, the energy spectral index is also fairly dependent on the inverse energy from the thermalized downstream region into the precursor region."," Additionally, the energy spectral index is also fairly dependent on the inverse energy from the thermalized downstream region into the precursor region."30" 1n summary, we performed the dynamical Monte Carlo simulations using the Gaussian scattering angular distributions based on the Matlab platform by monitoring the particle's mass, momentum and energy at any instant in time."," In summary, we performed the dynamical Monte Carlo simulations using the Gaussian scattering angular distributions based on the Matlab platform by monitoring the particle's mass, momentum and energy at any instant in time."31" The specific mass, momentum and energy loss functions with respect to time are presented."," The specific mass, momentum and energy loss functions with respect to time are presented."32 A series of analyses of the particle losses are obtained in the four cases., A series of analyses of the particle losses are obtained in the four cases.33" We successfully examine the relationship between the shock compression ratio and the energy losses, as well as verify the consistency of the energy spectral index with the inverse energy injected from the downstream to precursor region in the simulation cases which are applied with the preseribed Gaussian scattering angular distributions."," We successfully examine the relationship between the shock compression ratio and the energy losses, as well as verify the consistency of the energy spectral index with the inverse energy injected from the downstream to precursor region in the simulation cases which are applied with the prescribed Gaussian scattering angular distributions."34" In conclusion, the relationship of the shock compression ratio with the energy losses via FEB verify that the energy spectral index is determined by the inverse energy function with time."," In conclusion, the relationship of the shock compression ratio with the energy losses via FEB verify that the energy spectral index is determined by the inverse energy function with time."35" In fact, these energy losses simultaneously depend on the assumption of the prescribed scattering law."," In fact, these energy losses simultaneously depend on the assumption of the prescribed scattering law."36" As expected, the maximum energy of accelerated particles is limited by the size of the FEB according to the maximum mean free path in each case."," As expected, the maximum energy of accelerated particles is limited by the size of the FEB according to the maximum mean free path in each case."37" However, there is still a fairly large difference between the maximum energy of the particle from the different cases with the same size of the FEB."," However, there is still a fairly large difference between the maximum energy of the particle from the different cases with the same size of the FEB."38" We find that the total energy spectral index increases as the standard deviation value of the scattering angular distribution increases, but the subshock’s energy spectral index decreases as the standard deviation value of the scattering angular distribution increases."," We find that the total energy spectral index increases as the standard deviation value of the scattering angular distribution increases, but the subshock's energy spectral index decreases as the standard deviation value of the scattering angular distribution increases."39" In these multiple scattering angular distribution simulations, the prescribed scattering law dominates the energy losses and the inverse energy."," In these multiple scattering angular distribution simulations, the prescribed scattering law dominates the energy losses and the inverse energy."40" Consequently, the case of applying a preseribed law which leads to the minimum energy losses will produce a harder subshock's energy spectrum than those in the cases with larger energy losses."," Consequently, the case of applying a prescribed law which leads to the minimum energy losses will produce a harder subshock's energy spectrum than those in the cases with larger energy losses."41" These relationships will drive us to find a newly prescribed scattering law which leads to the minimum energy losses, making the shock compression ratio more closely approximate the standard value of four for a nonrelativistic shock with high Mach number in astrophysics."," These relationships will drive us to find a newly prescribed scattering law which leads to the minimum energy losses, making the shock compression ratio more closely approximate the standard value of four for a nonrelativistic shock with high Mach number in astrophysics."42Values of Tr. logg. £ and individual abundances estimated in this way were then used as initial guesses for starting another iterative procedure based. on. the ATLASI2 code.,"Values of $_{\rm eff}$ , $\log g$, $\xi$ and individual abundances estimated in this way were then used as initial guesses for starting another iterative procedure based on the ATLAS12 code."43 The best fit was obtained after three —iterations and led to the following parameters: Fr 17400 + 150 Wk. g log44.0 + 0.1. £——44.200.3 km s and οsin? 220.5+00.5 km ," The best fit was obtained after three iterations and led to the following parameters: $_{\rm eff}$ 7400 $\pm$ 150 K, $\log g$ 4.0 $\pm$ 0.1, $\xi$ $\pm$ 0.3 km $^{-1}$ and $v_e \sin i$ $\pm$ 0.5 km $^{-1}$ ."44The corresponding abundances are shown in the second. column of Table 2.. ltvabchikova., The corresponding abundances are shown in the second column of Table \ref{abund}. .45IWochukhoy&Bagnulo(2008) analvzed 11D227411 in a study of Xp. stars. and derived.— the following parameters: “Pay 77650 In. logg 44.0. οΜΗ 118.5 kms toandé==2?2.5 kms +.," \citet{ryab08}46 analyzed 27411 in a study of Ap stars and derived the following parameters: $_{\rm eff}$ 7650 K, $\log g$ 4.0, $v_e \sin i$ 18.5 km $^{-1}$, and $\xi$ 2.5 km $^{-1}$."47 Considering the experimental errors. these values are in agreement with Ours.," Considering the experimental errors, these values are in agreement with ours."48 The fits between the observed. anclsynthetic Balmer lines are shown in Fig. Ll. , The fits between the observed andsynthetic Balmer lines are shown in Fig. \ref{balmer}. .49Phe determination of surface gravity was constrained by using the triplet at AA 5183 Alas shown in Fig. 2..," The determination of surface gravity was constrained by using the triplet at $\lambda\lambda$ 5167--5183 , as shown in Fig. \ref{MgIb}. ."50 Errors in Ty and logg were estimated by the change in parameter values which leads to an increase of us by unity (Lamptonetal.1976)., Errors in $_{\rm eff}$ and $\log g$ were estimated by the change in parameter values which leads to an increase of $\chi^2$ by unity \citep{lampton76}.51. Η we adopt Zi;=74003:150 Ix and logg=4.00£0.10 from our spectroscopic analysis. we may use the relationships by 'lorresctal.(2010). to derive logL/h.=0.99+ 0.12.," If we adopt $T_{\rm eff} = 7400 \pm 150$ K and $\log g = 4.00 \pm 0.10$ from our spectroscopic analysis, we may use the relationships by \citet{torres10}52 to derive $\log L/L_\odot = 0.99 \pm 0.12$ ."53 These relate the mass and radius of a star to the elfective temperature ancl eravity through empirical calibrations., These relate the mass and radius of a star to the effective temperature and gravity through empirical calibrations.54 The greatest. source of uncertainty is the surface. gravity determination., The greatest source of uncertainty is the surface gravity determination.55 TheHipparcos parallax [or 227411. 4 -—-—]11.1343 00.38 (VanLeeuwen2007).. is useful in refining the location of the star in the Lk diagram.," The parallax for 27411, $\pi$ $\pm$ 0.38 \citep{van07}, is useful in refining the location of the star in the HR diagram."56 We show below that some caution is required. in estimating reddening in Am/Pim stars and that 227411. is. not significantlyreddened., We show below that some caution is required in estimating reddening in Am/Fm stars and that 27411 is not significantlyreddened.57 We therefore adopt. το=6.075. which gives an absolute magnitude A4;=1.3140.08 where the error is derived. from the error in the parallax.," We therefore adopt $V_0 = 6.075$, which gives an absolute magnitude $M_V = 1.31 \pm 0.08$ where the error is derived from the error in the parallax."58 If we adopt the bolometric correction DC = 0.051 derived rom Balona(1994).. we have Ad.)=1.36x0.09.," If we adopt the bolometric correction BC = 0.051 derived from \citet{balona94}, we have $M_{\rm bol} = 1.36 \pm 0.09$."59 Using Mi. ==44.74 (Drilling&Landolt1999).. we obtain og(L/L.)=1.3540.07.," Using $_{\rm bol,\odot}$ 4.74 \citep{drilling99}, we obtain $\log(L/L_\odot) = 1.35 \pm 0.07$."60 From the luminosity ancl using dng 77400 WK. we obtain 2/42.=2.8840.10.," From the luminosity and using $T_{\rm eff}$ 7400 K, we obtain $R/R_\odot = 2.88 \pm 0.10$."61 Phe surface gravity obtainedby using the parameters derived from the xwallax and assuminge a massof about 2.17£0.05AL. is oggBs73.9d 0.1., The surface gravity obtainedby using the parameters derived from the parallax and assuming a massof about $2.17 \pm 0.05~M_\odot$ is $\log g \approx 3.9 \pm 0.1$ .62" Allthe astrophysical quantitiesi derivedrere are summarizedin ""Fable 3..", Allthe astrophysical quantities derivedhere are summarizedin Table \ref{phot}. .63 The location of the star in the II diagram. togetherwithsome evolutionary tracks computed. for. non-solar," The location of the star in the HR diagram, togetherwithsome evolutionary tracks computed for non-solar"64We would like to thank Hans-Walter Rix for clarifying discussions and the organizers of the Sixth Harvard-Smithsonian Sackler Conference. where this work was completed.,"We would like to thank Hans-Walter Rix for clarifying discussions and the organizers of the Sixth Harvard-Smithsonian Sackler Conference, where this work was completed."65 OG is supported in part by NSF grant AST-, OG is supported in part by NSF grant AST-0708087.66of our analvsis using the same model atmosphere and non-LTE line formation codes ancl (he same technique in fitting Balmer line profiles is essential for the internal accuracy of our method.,of our analysis using the same model atmosphere and non-LTE line formation codes and the same technique in fitting Balmer line profiles is essential for the internal accuracy of our method.67 This is why we do not include the results of other published quantitative studies of additional objects. for which we do not have spectra al our disposal.," This is why we do not include the results of other published quantitative studies of additional objects, for which we do not have spectra at our disposal."68 Fig., Fig.69 3 shows the FGLR. with all objects included., 3 shows the FGLR with all objects included.70 The regression coefficients are very similar (d=—3. 11. b= 13.49) and the standard. deviation of the residual bolometric magnitude has only slightly increased (ο σ = 0.28 mag.," The regression coefficients are very similar $a=-3.71$ , $b=13.49$ ) and the standard deviation of the residual bolometric magnitude has only slightly increased to $\sigma$ = 0.28 mag."71 We note that the scatter increases αἱ absolute magnitudes brighter than —3 mag (see discussion below)., We note that the scatter increases at absolute magnitudes brighter than $-8$ mag (see discussion below).72 Fie., Fig.73 3 also includes the relationship obtained from the evolutionary calculations including stellar rotation and mass-loss (Alevuet&Maecder 2000))., 3 also includes the relationship obtained from the evolutionary calculations including stellar rotation and mass-loss \citealt{meynet00}) ).74 While the slopes practically agree (note that the stellar evolution relationship shows a slight curvature reflecting the mild change in (he mass-Iuninosity exponent a). there is a small offset in log(g ) bv 0.07 dex.," While the slopes practically agree (note that the stellar evolution relationship shows a slight curvature reflecting the mild change in the mass-luminosity exponent $\alpha$ ), there is a small off-set in $g$ ) by 0.07 dex."75 This can be the result of à svstematic elect in the determination of gravity or in the temperature scale. but it can also indicate a small deficiency of the evolutionary models.," This can be the result of a systematic effect in the determination of gravity or in the temperature scale, but it can also indicate a small deficiency of the evolutionary models."76 In aly case. as long as an empirical calibration of the relationship with stus at known distances is used. (he accuracy of the distance determination will not be affected.," In any case, as long as an empirical calibration of the relationship with stars at known distances is used, the accuracy of the distance determination will not be affected."77 The results presented in (he previous sections are verv encouraging., The results presented in the previous sections are very encouraging.78 The flux-weishted eravities of late D and early A supergiants awe obviously very tightly correlated with absolute bolometrie magnitude., The flux-weighted gravities of late B and early A supergiants are obviously very tightly correlated with absolute bolometric magnitude.79 The application of this relationship. once properly calibrated. [or extragalactic distance determinations is streüghtlorward.," The application of this relationship, once properly calibrated, for extragalactic distance determinations is straightforward."80 It requires multi-colour photometry of galaxies containing a voung stellar population to identify possible blue supergiants and subsequent medium resolution (5 A)) multi-objeet spectroscopy (see 20023) to determine effective temperature and eravity directly [rom the spectra.," It requires multi-colour photometry of galaxies containing a young stellar population to identify possible blue supergiants and subsequent medium resolution $\sim 5\,$ ) multi-object spectroscopy (see \citealt{bresolin01,bresolin02}) ) to determine effective temperature and gravity directly from the spectra."81 The spectral analvsis will also vield bolometric correction (which is small for these spectral tvpes) and intrinsic colour so Chat an accurate correction [or reddening and extinction is possible., The spectral analysis will also yield bolometric correction (which is small for these spectral types) and intrinsic colour so that an accurate correction for reddening and extinction is possible.82 Application of the FGLR. will then provide the absolute bolometric magnitude. which by comparison with the de-reddened visual magnitude will eive the distance modulus.," Application of the FGLR will then provide the absolute bolometric magnitude, which by comparison with the de-reddened visual magnitude will give the distance modulus."83 Assuming a residual scatter of σ = 0.3 mag lor the FGLI (see Fig., Assuming a residual scatter of $\sigma$ = 0.3 mag for the FGLR (see Fig.84 2 and 3) we estimate that with 10 supereiant stars per galaxy we can achieve an accuracy oL 0.1 mag in distance modulus., 2 and 3) we estimate that with 10 supergiant stars per galaxy we can achieve an accuracy of $0.1$ mag in distance modulus.85 We are confident that in one night of observing time we can reach down to V.=22.5 with the existing verv efficient. medium resolution multi-object spectrographs attached. to 81-class telescopes., We are confident that in one night of observing time we can reach down to $V = 22.5$ with the existing very efficient medium resolution multi-object spectrographs attached to 8m-class telescopes.86 With objects in an absolute magnitude range between. —5 and —10 mae theFGLHB method appears to be applicable out to distance moduli of ο—AL=30.5 or even, With objects in an absolute magnitude range between $-8$ and $-10$ mag theFGLR method appears to be applicable out to distance moduli of $m-M = 30.5$ or even87source {5 offset from the cluster with a confidence of 3 sigma.,source is offset from the cluster with a confidence of 3 sigma.88 We have also identitied a new ULX. CXOU J230453.0+121959. that is associated with a young massive supercluster of a mass ~10°M.. ," We have also identified a new ULX, CXOU J230453.0+121959, that is associated with a young massive supercluster of a mass $\sim10^{6}$ "89as observed in the AS cliunel. our analysis suggests for instance au uncertaimtv for the mass as laree as if ouly 20 modes are cousidered (for the closest models. effective temperatures are compatible within GOR. and radi and DIuniuosities differ bv ~ A).,"as observed in the AS channel, our analysis suggests for instance an uncertainty for the mass as large as if only 20 modes are considered (for the closest models, effective temperatures are compatible within 60K, and radii and luminosities differ by $\sim$ )."90 With the method based on Oxia ALL uncertantv of for the mass rendus Oy up to 50 considered modes (for the closest models. effective temperatures are conipatible within GOI. radii differ bv and Iundnuosities are equal).," With the method based on $ \delta_{sls}$, an uncertainty of for the mass remains for up to 50 considered modes (for the closest models, effective temperatures are compatible within 60K, radii differ by and luminosities are equal)."91of the amplitude of the intrinsic spectrum. and of enhanced soft X-ray absorption around the minimum flux of the SO cycle.,"of the amplitude of the intrinsic spectrum, and of enhanced soft X-ray absorption around the minimum flux of the SO cycle."92 These features are fully compatible with the model of a bulge around the centre of a precessing disc. proposed by ΡΖΙ05.," These features are fully compatible with the model of a bulge around the centre of a precessing disc, proposed by PZI08."93 We note that the precession mechanism should be compatible with the change of its period by a factor of 2., We note that the precession mechanism should be compatible with the change of its period by a factor of 2.94 This appears easier ο achieve with radiation driven precession than with that caused surely by the tidal forces exerted by the donor (e.g.. Wijers&Pringle 1999:: ΟΡΟΙ: Capronietal. 20065.," This appears easier to achieve with radiation driven precession than with that caused purely by the tidal forces exerted by the donor (e.g., \citealt{wp99}; OD01; \citealt{caproni06}) )."95" ODOL tind that Cyg X-l| lies in a region of marginal stability of radiation induced orecession. Which possibly ean explain the observed large change of Pi, as a manifestation of a chaotic behaviour."," OD01 find that Cyg X-1 lies in a region of marginal stability of radiation induced precession, which possibly can explain the observed large change of $\psup$ as a manifestation of a chaotic behaviour."96 Still. if this is the case. the rather coherent ~ 150-d and ~300-d SO modulations over ime scales of years are difficult to understand.," Still, if this is the case, the rather coherent $\sim$ 150-d and $\sim$ 300-d SO modulations over time scales of years are difficult to understand."97 An alternative is that the SO modulation is not chaotic but stable with the period of ~300 d. but there are changes in the nodal precession mode.," An alternative is that the SO modulation is not chaotic but stable with the period of $\sim$ 300 d, but there are changes in the nodal precession mode."98 The observed period of =300 d would then correspond to the main. retrograde mode with v7=0 bending nodes CODOIL.," The observed period of $\simeq 300$ d would then correspond to the main, retrograde mode with $m=0$ bending nodes (OD01)."99 The «150-d observed period would then correspond to the prograde 7j=| mode (ODOL). in which the actual precession period would be still ~300 d. but the disc bending gives rise to two maxima/minima of the observed flux per period.," The $\sim$ 150-d observed period would then correspond to the prograde $m=1$ mode (OD01), in which the actual precession period would be still $\sim$ 300 d, but the disc bending gives rise to two maxima/minima of the observed flux per period."100 An argument in favour of this interpretation is that the precession during the ~ 150- modulation was found prograde by PZIOS., An argument in favour of this interpretation is that the precession during the $\sim$ 150-d modulation was found prograde by PZI08.101 In this model. the absorption periodically varving with a half of the precession period needs to be due to some medium around the bent dise.," In this model, the absorption periodically varying with a half of the precession period needs to be due to some medium around the bent disc."102" We have found that the ASM count rates averaged over the hard states characterized by Pi,=150 d and Py)=300 d are almost identical.", We have found that the ASM count rates averaged over the hard states characterized by $\psup\simeq 150$ d and $\psup\simeq 300$ d are almost identical.103 The bolometric fluxes in these two epoch are also found to be almost identical (Zdziarskietal.20113..., The bolometric fluxes in these two epoch are also found to be almost identical \citep{z11}.104 Thus. the change of the SO period is not associated with a significant change of the average luminosity (most likely tied to the accretion rate). which appears to present a difficulty for the model of driven precession.," Thus, the change of the SO period is not associated with a significant change of the average luminosity (most likely tied to the accretion rate), which appears to present a difficulty for the model of radiation-driven precession."105 We also study the orbital modulation of the X-ray light curves., We also study the orbital modulation of the X-ray light curves.106 We find that the depth of this modulation in the soft X-rays is significantly higher during 2006—2010 than before., We find that the depth of this modulation in the soft X-rays is significantly higher during 2006–2010 than before.107 This also manifests itself in the frequency of the X-ray dips. much higher in this epoch than before.," This also manifests itself in the frequency of the X-ray dips, much higher in this epoch than before."108 This appears to indicate that the stellar wind. for some reasons. has become more anisotropic than before.," This appears to indicate that the stellar wind, for some reasons, has become more anisotropic than before."109 Likely. the bulge at the dise centre may have now a higher optical depth than before.," Likely, the bulge at the disc centre may have now a higher optical depth than before."110 We tind statistically significant orbital modulation of the flux (but not of the spectral slope) in the BAT light curve., We find statistically significant orbital modulation of the flux (but not of the spectral slope) in the BAT light curve.111 The cause of the modulation appears to be Compton scattering of photons away from the line of sight (LO6)., The cause of the modulation appears to be Compton scattering of photons away from the line of sight (L06).112 As the optical depth of the wind is higher at the superior conjunction than at the inferior one. this scattering causes changes of the observed hard X-ray flux by about £2 per cent.," As the optical depth of the wind is higher at the superior conjunction than at the inferior one, this scattering causes changes of the observed hard X-ray flux by about $\pm 2$ per cent."113 Finally. we find no correlation of the X-ray and radio properties of the source with the reported detections in high energy y--rays.," Finally, we find no correlation of the X-ray and radio properties of the source with the reported detections in high energy -rays."114 The TeV detection happened during a relatively high X-ray flux (see also Malzaeetal. 2008)). whereas the GeV ones took place when both the X-ray and radio fluxes were low.," The TeV detection happened during a relatively high X-ray flux (see also \citealt{malzac08}) ), whereas the GeV ones took place when both the X-ray and radio fluxes were low."115 We note that LAT detected no GeV emission from Cyg X-I (Hilletal.2010) during the first two flares reported by (Sabatinietal.2010a:Bulgarelli2010).," We note that LAT detected no GeV emission from Cyg X-1 \citep{hill10} during the first two flares reported by \citep{sabatini10a,bulgarelli10}."116 We thank A. Ogorzalek and L. Wen for help with calculating the rescaled periodograms. H. Krimm and P. Lubinsski for help with the BAT data. G. Dubus for discussion of radiation induced precession. and the referee for valuable suggestions.," We thank A. ek and L. Wen for help with calculating the rescaled periodograms, H. Krimm and P. Lubińsski for help with the BAT data, G. Dubus for discussion of radiation induced precession, and the referee for valuable suggestions."117 This research has been supported in part by the Polish MNiSW grants NN203065933 and. 362/1/N-INTEGRAL/2008/09/0., This research has been supported in part by the Polish MNiSW grants NN203065933 and 362/1/N-INTEGRAL/2008/09/0.118 The AMI Arrays are operated by the University of Cambridge and supported by the STFC., The AMI Arrays are operated by the University of Cambridge and supported by the STFC.119 We acknowledge the use of data obtained through the HEASARC online service provided by NASA/GSFC., We acknowledge the use of data obtained through the HEASARC online service provided by NASA/GSFC.120eutron stars have the strongest magnetic fields found in the universe. with fields perhaps as large as 107 G for so-called magnetars MMüuralkumi 1999). around LOL C for voung (~10 vear) radio and X-ray. pulsars. and a still appreciable 10.1027 € for much older (~102 vear) millisecond pulsars CChanmugam 1992: Bhattacharva 1995: Lvne 2000).,"Neutron stars have the strongest magnetic fields found in the universe, with fields perhaps as large as $10^{15}$ G for so-called magnetars Murakami 1999), around $10^{12}$ G for young $\sim10^7$ year) radio and X-ray pulsars, and a still appreciable $10^8 - 10^{10}$ G for much older $\sim10^{10}$ year) millisecond pulsars Chanmugam 1992; Bhattacharya 1995; Lyne 2000)."121 ‘This correlation between field strength and age suggests that jese very dilferent strengths are due to the field decaving in time. rather than to any intrinsic. dilferences. between lilferent neutron stars.," This correlation between field strength and age suggests that these very different strengths are due to the field decaying in time, rather than to any intrinsic differences between different neutron stars."122 One would therefore like to identify 10 processes causing the field to decay., One would therefore like to identify the processes causing the field to decay.123 The additional observation that most weakly magnetic =reutron stars have binary companions. whereas very [ew strongly magnetic ones do BBhattacharva 1995). suggests that aceretion of mass from the companion is somehow causing the field to decay (by mechanisms that need not concern us here. but see for example. Blonclin Freese 1986: Romani 1990: Urpin Geppert 1995).," The additional observation that most weakly magnetic neutron stars have binary companions, whereas very few strongly magnetic ones do Bhattacharya 1995), suggests that accretion of mass from the companion is somehow causing the field to decay (by mechanisms that need not concern us here, but see for example Blondin Freese 1986; Romani 1990; Urpin Geppert 1995)."124 The observational evidence is unfortunately inconclusive. with Taam van den LHeuvel (1986) claiming a correlation between field strength and accreted mass. but Wijers (1997) disputing this.," The observational evidence is unfortunately inconclusive, with Taam van den Heuvel (1986) claiming a correlation between field strength and accreted mass, but Wijers (1997) disputing this."125 One would therefore like to consider the possibility of other mechanisms besides accretion., One would therefore like to consider the possibility of other mechanisms besides accretion.126 One such alternative is Hall drift. first. proposed. by Jones (1988). in which the magnetic field. influences itself. through a quadratic nonlinearity.," One such alternative is Hall drift, first proposed by Jones (1988), in which the magnetic field influences itself through a quadratic nonlinearity."127 Lit is relevant at all. Hall cirift will therefore be most important for the very strongest fields which as we saw Conc to occur in isolated neutron stars. where accretion is not acting at all.," If it is relevant at all, Hall drift will therefore be most important for the very strongest fields – which as we saw tend to occur in isolated neutron stars, where accretion is not acting at all."128 Llall απ is thus likely to be thedominant mechanism inlluencing the magnetic fields of these stars., Hall drift is thus likely to be the dominant mechanism influencing the magnetic fields of these stars.129 Of course. it could potentially be important in binaries as well. at least in the early stages while their fields are still relatively strong.," Of course, it could potentially be important in binaries as well, at least in the early stages while their fields are still relatively strong."130 Again as a result of this quadratic nonlinearity. the timescale on which Llall drift. might be expected. to act is almost necessarily inversely proportional to |B].," Again as a result of this quadratic nonlinearity, the timescale on which Hall drift might be expected to act is almost necessarily inversely proportional to $|{\bf B}|$."131 Jones sugeests that it is given by, Jones suggests that it is given by132We will here consider an alternative method to determine a dipole signature which tries to overcome some of the difficulties just mentioned aud apply it to simulated data sets.,We will here consider an alternative method to determine a dipole signature which tries to overcome some of the difficulties just mentioned and apply it to simulated data sets.133 Let d be the dipole’s direction. aud 5=acosto+d) the angle between the dipole and the eveuts arrival direction à.," Let $\hat d$ be the dipole's direction, and $\gamma={\rm acos}({\hat134 n}\cdot \hat d)$ the angle between the dipole and the event's arrival direction ${\hat n}$."135 A dipolar distribution should give rise toa CR fiux (incident ou the Earth) of the form where 6 is the dipole’s amplitude (1.0. D=ad)., A dipolar distribution should give rise to a CR flux (incident on the Earth) of the form where $a$ is the dipole's amplitude (i.e. $\vec D=a\hat d$ ).136 When this flux is observed by an experiment with non-uniform exposure Gutcerated over time) wr). the expected event rates should behave as with the normalization. which depends ou α aud d. fixed to reproduce the total umber of events observed.," When this flux is observed by an experiment with non-uniform exposure (integrated over time) $\omega(\hat n)$, the expected event rates should behave as with the normalization, which depends on $a$ and $\hat d$, fixed to reproduce the total number of events observed."137" Suppose that now we try to fit a dipole signal along a direction d. along which the expected distribution would bewhere 0 is the azimthal auele around the axis d"". cos5—d«9 aud where cos}—dνα aud 05 is the azimuthal angle. mecasured around Zl. of he dipole vector."," Suppose that now we try to fit a dipole signal along a direction $\hat d'$ , along which the expected distribution would bewhere $\theta'$ is the azimuthal angle around the axis $\hat d'$, $\cos\gamma'\equiv \hat d'\cdot \hat n$ and where $\cos\beta\equiv \hat d\cdot \hat d'$ and $\theta_d'$ is the azimuthal angle, measured around $\hat d'$, of the dipole vector."138 It is then clear that if; were uniform m the sky. the terii proportional o cos?—(7) in the above integral wouldvanish. leading to the behavior (ANdeos!x(1|ecoscoss). aud hence the dipole amplitude inferred. xuriug statistical fluctuations. would be ecos. which is just thedipole componoeut along the d’ axis.," It is then clear that if $\omega$ were uniform in the sky, the term proportional to $\cos(\theta'-\theta_d')$ in the above integral wouldvanish, leading to the behavior $N/{\rm d}\cos\gamma'\propto(1+a\cos\beta\cos\gamma')$, and hence the dipole's amplitude inferred, barring statistical fluctuations, would be $a\cos\beta$, which is just the dipole component along the $\hat d'$ axis."139 One may then envisage that the dipole’s direction could © obtained as the oue maximizing the recoustructed dipole's amplitude. but rowever. for non-unifoin exposures the cos(0—05) tei does not average to zero in eeucral. so that this procedure could lead to a biased result.," One may then envisage that the dipole's direction could be obtained as the one maximizing the reconstructed dipole's amplitude, but however, for non-uniform exposures the $\cos(\theta'-\theta_d')$ term does not average to zero in general, so that this procedure could lead to a biased result."140 There is however a particularly relevant case where this bias is absent. which is when one considers d2 τν ie. the NS equatorial axis. since for the case of uifon exposure iu right ascension (with au arbitrary declination dependence). the cos(0—05) term in the integral willindeed. vauislil.. hence leadiug to a behavior dNdeos5x(1|04co8 5). with a.=osinó being theamplitude of," There is however a particularly relevant case where this bias is absent, which is when one considers $\hat d'=\hat z$ , i.e. the NS equatorial axis, since for the case of uniform exposure in right ascension (with an arbitrary declination dependence), the $\cos(\theta'-\theta_d')$ term in the integral willindeed , hence leading to a behavior $N/{\rm d}\cos\gamma'\propto(1+a_z\cos\gamma')$ , with $a_z=a\sin\delta$ being theamplitude of"141If accuracy is required. then a µας survey is likely to measure the mass ancl orbital parameters of one planet with a mass 7 or 2AL lor five- or teu-vear surveys. respectively.,"If accuracy is required, then a 1 $\mu$ as survey is likely to measure the mass and orbital parameters of one planet with a mass $\sim7$ or $2 M_\oplus$ for five- or ten-year surveys, respectively."142 Süunilarly. a 2 µας survey is likely to measure with accuracy the mass aid orbital parameters of one planet with a mass 5 or BAL for five- or ten-year surveys. respectively (see £1).," Similarly, a 2 $\mu$ as survey is likely to measure with accuracy the mass and orbital parameters of one planet with a mass $\sim8$ or $3 M_\oplus$ for five- or ten-year surveys, respectively (see 4)."143 The two-tier strategy would likely measure the mass and orbital parameters with accuracy for a SAL planet in five vears aud a ~3M. planet in teu years (see 110)., The two-tier strategy would likely measure the mass and orbital parameters with accuracy for a $\sim8 M_\oplus$ planet in five years and a $\sim3 M_\oplus$ planet in ten years (see 10).144 We can estimate the number of planets of all masses that SIM is likely to find., We can estimate the number of planets of all masses that SIM is likely to find.145 These estimates are much more reliable than the previous ones. since radial velocity surveys are sensitive to the vpical planets that SIM is expected to fiuc.," These estimates are much more reliable than the previous ones, since radial velocity surveys are sensitive to the typical planets that SIM is expected to find."146 Our estimates [or a five-year nission duratiou are also uore reliable thau for a ten-year mission. since radial velocity surveys are just beginuiug to detect jxanets with orbital periods 10 years.," Our estimates for a five-year mission duration are also more reliable than for a ten-year mission, since radial velocity surveys are just beginning to detect planets with orbital periods $\sim10$ years."147 If its primary goal is to detect and measure masses aud orbits lor the greatest. number of janets. regardless of mass. then SIN. should pursue a survey of a large number of stars with a 'elatively low single measurement accuracy.," If its primary goal is to detect and measure masses and orbits for the greatest number of planets, regardless of mass, then SIM should pursue a survey of a large number of stars with a relatively low single measurement accuracy."148 For example. using a siugle measurement precision of lor 2 µας. SIM would be expected to detect 27+Y or 69+13 planets in a five-vear mission or ~3348 or 9815 planets iu a teu-vear mission (see £1).," For example, using a single measurement precision of 1 or 2 $\mu$ as, SIM would be expected to detect $\sim27\pm7$ or $69\pm13$ planets in a five-year mission or $\sim33\pm8$ or $98\pm15$ planets in a ten-year mission (see 4)."149 The two-tier strategy that we consider would likely discover ~108£21 or 162421 planets for a five- or ten-vear mission. respectively (see 99).," The two-tier strategy that we consider would likely discover $\sim108\pm21$ or $162\pm24$ planets for a five- or ten-year mission, respectively (see 9)."150 Racial velocity surveys are also capable of discovering giant planets with orbital periods of several vears., Radial velocity surveys are also capable of discovering giant planets with orbital periods of several years.151" Thus. it is important that we estimate the number of ""uew planets which SIM will detect that would not be discovered by a radial velocity survey of the same stars."," Thus, it is important that we estimate the number of “new” planets which SIM will detect that would not be discovered by a radial velocity survey of the same stars."152 A ten-year racial velocity survey with 3 precision would be expected to detect 8+ land 31-9 planets. lor the 1 and 2 µας target lists.," A ten-year radial velocity survey with $3$ precision would be expected to detect $\sim8\pm4$ and $34\pm9$ planets, for the 1 and 2 $\mu$ as target lists."153 For the same two target lists. SIM would be expected to detect 17 or LE new planets ini a five-year mission aud ~25 or GS new planets in a teu-vear mission (see £1).," For the same two target lists, SIM would be expected to detect $\sim17$ or $44$ new planets in a five-year mission and $\sim25$ or $68$ new planets in a ten-year mission (see 4)."154 The two-tier strategy would be expected to detect ~57 or OF new planets lor a five- or ten-year uiission., The two-tier strategy would be expected to detect $\sim57$ or $97$ new planets for a five- or ten-year mission.155 However. conducting such a large radial velocity survey and includiug a siguilicaut Craction ol M stars would require a large amount of observing time.," However, conducting such a large radial velocity survey and including a significant fraction of M stars would require a large amount of observing time."156 Unfortunately. SIM will not obtain accurate planetary masses or orbits for mauy of the planets it detects.," Unfortunately, SIM will not obtain accurate planetary masses or orbits for many of the planets it detects."157 For example. using a single measurement precision of 1 or 2 gras. SIM would be expected to determine masses aud orbits with accuracy for 1646 or LL11 planets for a five-year inission or 22+8 or 66+12 planets for a ten-year mission (~65% of the planets detected with the same mission parameters).," For example, using a single measurement precision of 1 or 2 $\mu$ as, SIM would be expected to determine masses and orbits with accuracy for $\sim16\pm6$ or $44\pm11$ planets for a five-year mission or $\sim22\pm8$ or $66\pm12$ planets for a ten-year mission $\sim65\%$ of the planets detected with the same mission parameters)."158 SIM. would be expected to determine masses and orbits with, SIM would be expected to determine masses and orbits with159The amplitudes of solax-like oscillatiois are defined the valance between excitation aud camping.,The amplitudes of solar-like oscillations are defined by the balance between excitation and damping.160 The oscillationsby are excited stochastically by the acoustic noise generated w the turbulent motion of the couvective elements., The oscillations are excited stochastically by the acoustic noise generated by the turbulent motion of the convective elements.161 Iu Samadi Goupil (2001. II hereatter)) a theoretical ornmulation for the acoustic power injected mto solar-like oscillations is proposed aud which supplements previous heories.," In Samadi Goupil (2001, I \nocite{Samadi00I}) ) a theoretical formulation for the acoustic power injected into solar-like oscillations is proposed and which supplements previous theories."162 We refer the reader to Samadi(2001) for a detailed summary and discussion on some recent unsolved xoblenms., We refer the reader to \cite*{Samadi01} for a detailed summary and discussion on some recent unsolved problems.163 The excitation process depeuds ou the asstuned spectrmm. as discussed. forexample. by Samadietal.(2001b) and Samadi(2001).," The excitation process depends on the assumed turbulence spectrum, as discussed, forexample, by \cite*{Samadi00b} and \cite*{Samadi01}."164. It also «epends crucially ou the convection model to compute the stratification of the convectively unstable lavers iu the equilibrium model., It also depends crucially on the convection model to compute the stratification of the convectively unstable layers in the equilibrium model.165 The αλλοι of cherev injected iuto the oscillations depends stronev on the velocity of the convective elements., The amount of energy injected into the oscillations depends strongly on the velocity of the convective elements.166 The main goal of tus work is fo asses changes in the oscillaion power spectrum due to modifications of the convection treatment in the equilibrium mool., The main goal of this work is to asses changes in the oscillation power spectrum due to modifications of the convection treatment in the equilibrium model.167 We consider two differeut formulations: the classical description of leugth theory bv DóhurViteuse(1958) aid a nonlocal eeneralizatiou of the mixing-lenueth formulation by Cough (1976. 1977)).," We consider two different formulations: the classical description of mixing-length theory by \cite*{Bohm58} and a nonlocal generalization of the mixing-length formulation by Gough (1976, \nocite{Gough76,Gough77}) )."168 Additiowilly we study the dependence of he oscillation power on the asstuned turbulence spectrum in the excitation model for both convection formulations., Additionally we study the dependence of the oscillation power on the assumed turbulence spectrum in the excitation model for both convection formulations.169" The.à theoreticalwrote: formulation""yati iu1 PaperI involves:Mx,Tau two-, Fee paraneters,", The theoretical formulation in I involves two free parameters.170 These parauieters are calilrated. SUCH as ον reproduce for a solar model the observed acoustic power PC:(ctefr (Sectionwavesel H2)., These parameters are calibrated such as to reproduce for a solar model the observed acoustic power spectrum (Section 2).171 In Sectionoecti 3D weaura computeH oscillationlati POW spectra for several stellar models using the two onvectiou91enu forformationatious., In Section 3 we compute oscillation power spectra for several stellar models using the two convection formulations.172 We couchide that Eddiugtons. perf.yu will. awv us to distiuguisli between the two treatments of convection and consequently that Eddinetou will provide further constraiu on convection models., We conclude that Eddington's performance will allow us to distinguish between the two treatments of convection and consequently that Eddington will provide further constraints on convection models.173 The acoustic power P injected into the oscillatious is defined (e.g. Goldreichetal.199 D) iu terms of the mode damping rate 4. the oscillation meau-quare auplitude (47). the mode inertia £ aud oscillation frequency uw as: The mean-square auplitude is defined. by both the excitation by turbulent couvection aud by the damping process.," The acoustic power $P$ injected into the oscillations is defined $e.g.$, \cite{GMK94}) ) in terms of the mode damping rate $\eta$, the oscillation mean-square amplitude $\langle A^2 \rangle$, the mode inertia $I$ and oscillation frequency $\omega$ as: The mean-square amplitude is defined by both the excitation by turbulent convection and by the damping process."174" It cau be written as where & is the radial displaceimieut cigenfiuuction. p the deusitv. Ais tlemixing leneth. e the vertical rms velocity of the convective elements: T7,0) is a function which turbulencencludes the second derivative of ἐν. aud 2Gn) describes the ratio of the excitation by the entropy fluctuations to that by the Revuolds fhctuatious."," It can be written as where $\displaystyle{\xi_{\rm r}}$ is the radial displacement eigenfunction, $\rho$ the density, $\Lambda$ is themixing length, $w$ the vertical rms velocity of the convective elements; $\mathcal{F}^2(\xi_{\rm r},m)$ is a function which includes the second derivative of $\xi_{\rm r}$ , and $\mathcal{R}^2(m)$ describes the ratio of the excitation by the entropy fluctuations to that by the Reynolds fluctuations."175 The source. functio δαίω.1) aud Sete.) describe the coutributious from the Reynolds and eutropy fluctuations. respectively arisim from the simaller scales of the turbuleu cascade.," The source functions $\mathcal{S}_{\rm R}(\omega,m)$ and $\mathcal{S}_{\rm S}(\omega,m)$ describe the contributions from the Reynolds and entropy fluctuations, respectively, arising from the smaller scales of the turbulent cascade."176 Detailed expressions: forJ αυ CA75. οSpain). >Safa). We‘R= aud FF“) are given in II. The source fictions inchide the turbulent kinetic energy spectrmm £ (4). and the turbulentspectrum of the eutropy misxine-fluctuations £(k) whichcan be relaed to E(k) ly a simple expression (ο... Samadictal. 20012... III," Detailed expressions for $\left < A^2 \right >$ , $\mathcal{S}_{\rm R}(\omega,m)$, $\mathcal{S}_{\rm S}(\omega,m)$, $\mathcal{R}^2$ and $\mathcal{F}^2$ are given in I. The source functions include the turbulent kinetic energy spectrum $E(k)$ , and the turbulentspectrum of the entropy fluctuations $E_s(k)$ whichcan be related to $E(k)$ by a simple expression $e.g.,$ \cite{Samadi00II}, , II"177the available parameter space of cluster initial conditions leading to runaway growth.,the available parameter space of cluster initial conditions leading to runaway growth.178" They argue that this is because initial mass segregationdecreases the collision rate of stars in the core as, due to the increased average mass, the number density is decreased."," They argue that this is because initial mass segregation$decreases$ the collision rate of stars in the core as, due to the increased average mass, the number density is decreased."179" This is in line with earlier N-body simulations suggesting that runaway growth can occur only when the cluster is initially sufficiently collisional, (?),, in contrast to predictions based only on the core collapse time."," This is in line with earlier $N$ -body simulations suggesting that runaway growth can occur only when the cluster is initially sufficiently collisional, \citep{P04}, in contrast to predictions based only on the core collapse time."180 This result should be tested for larger N as the nature of the collisional runaway changes from being dominated by three-body binary formation at low N to single collisions at high N., This result should be tested for larger $N$ as the nature of the collisional runaway changes from being dominated by three-body binary formation at low $N$ to single-single collisions at high $N$.181 One of the main uncertainties in star cluster evolution lies in determining the true initial mass function (IMF)., One of the main uncertainties in star cluster evolution lies in determining the true initial mass function (IMF).182" Often it is assumed that the IMF is a standard power-law (or power-laws with different indices in different mass ranges) with no primordial radial variation in the cluster (e.g., ??7)."," Often it is assumed that the IMF is a standard power-law (or power-laws with different indices in different mass ranges) with no primordial radial variation in the cluster \citep[e.g.,183][]{Sal1959,MS1979,Kr2001}."184" Deviations from the standard IMFs are observed in many clusters both at the high and low-mass ends (e.g.,7?).."," Deviations from the standard IMFs are observed in many clusters both at the high and low-mass ends \citep[e.g., ][]{Elme2004}."185 In particular at the high mass end MFs are observed to be generally flatter compared to standard Salpeter power-law in young massive clusters like the Arches cluster (??)..," In particular at the high mass end MFs are observed to be generally flatter compared to standard Salpeter power-law in young massive clusters like the Arches cluster \citep{Stolte2002,kim06}."186" 'The numerical method that has been used here to investigate the dynamical evolution of star clusters is the Monte Carlo method, based on the classic work of ? and described in detail in ?,andreferences therein..."," The numerical method that has been used here to investigate the dynamical evolution of star clusters is the Monte Carlo method, based on the classic work of \citet{Henon} and described in detail in \citet[][and references therein]{FreRas07}."187" In Monte Carlo simulations, N, the total number of stars in the cluster is dependent on the initial half mass relaxation time in the cluster."," In Monte Carlo simulations, $N$, the total number of stars in the cluster is dependent on the initial half mass relaxation time in the cluster."188" For a Plummer sphere it is given by (?,eq.2.63),, where γε~0.01 is the coulomb logarithm and rp, is the half mass radius."," For a Plummer sphere it is given by \citep[eq. 2.63]{Spit87}, where $\gamma_c\sim0.01$ is the coulomb logarithm and $r_h$ is the half mass radius."189 Since our code now includes an explicit treatment of all stellar collisions we briefly summarize here the ‘sticky sphere’ method for stellar collisions (?).., Since our code now includes an explicit treatment of all stellar collisions we briefly summarize here the `sticky sphere' method for stellar collisions \citep{FreBen02}.190" In the sticky sphere approximation a collision occurs whenever the centers of two stars pass within a distance d=(R4+Re), with Ri,» being the stellar radii."," In the sticky sphere approximation a collision occurs whenever the centers of two stars pass within a distance $d= (R_{1}+R_{2})$, with $R_{1,2}$ being the stellar radii."191" Until this distance is reached, the gravitational influence of other stars as well as any mutual tidal interactions are neglected."," Until this distance is reached, the gravitational influence of other stars as well as any mutual tidal interactions are neglected."192" T'he cross section for such a collision is given by ?,section7.5.8, where Όπιας is the largest impact parameter leading to contact, Ure; is the relative velocity between two stars and v?=(2G(M,+M2)/(Ri+R3))95."," The cross section for such a collision is given by \citet[section 7.5.8]{Bin87}, where $b_{max}$ is the largest impact parameter leading to contact, $v_{rel}$ is the relative velocity between two stars and $v^{12}_{*} = (2G(M_{1}+M_{2})/(R_{1}+R_{2}))^{0.5}$ ."193" In a cluster where all stars have the same mass M. and radius R,,, the average local collision time Tyo is given by ?,section7.5.8 where n, is the stellar number density."," In a cluster where all stars have the same mass $M_{*}$ and radius $R_{*}$, the average local collision time $T_{coll}$ is given by \citet[section 7.5.8]{Bin87}194 where $n_{*}$ is the stellar number density."195 For a Maxwellian velocity distribution this becomes where oy is the velocity dispersion., For a Maxwellian velocity distribution this becomes where $\sigma_{v}$ is the velocity dispersion.196" In order to resolve collisional processes in a MC simulation we constrain the time step size δὲ according to an estimate of the central collision time using where f=5x107° is a constant chosen small enough to ensure that collisions are sampled sufficiently, and Τωμι is an estimate of Τερι based on equation 4,, and given by with quantities in angular brackets being local averages."," In order to resolve collisional processes in a MC simulation we constrain the time step size $\delta t$ according to an estimate of the central collision time using where $f= 5\times 10^{-3}$ is a constant chosen small enough to ensure that collisions are sampled sufficiently, and $\tilde{T}_{coll}$ is an estimate of $T_{coll}$ based on equation \ref{tcoll}, , and given by with quantities in angular brackets being local averages."197" The collision probability for two neighboring stars is calculated as where n, is a local estimate of the stellar number density.", The collision probability for two neighboring stars is calculated as where $n_{*}$ is a local estimate of the stellar number density.198" The time step size δὲ, is chosen"," The time step size $\delta t$ , is chosen"199iiomenutuni in units of οΑς. aud the nuubers are typical SAINS anodel values (Salgado et al.,"momentum in units of $G M^2/c$, and the numbers are typical SMNS model values (Salgado et al."200 1991. Cook. Shapiro Teukolkky 1991. VS98).," 1994, Cook, Shapiro Teukolsky 1994, VS98)."201 The spin-down time (i.c. the SN-CRB delay time) and corresponding wind luninositv can be estimated from the magnetic dipole formula (Pacini 1967. VS98) aud where HR.=LOR.και is a typical SAINS equatorial radius.," The spin-down time (i.e. the SN-GRB delay time) and corresponding wind luminosity can be estimated from the magnetic dipole formula (Pacini 1967, VS98) and where $R_* = 15 R_{*,15} {\rm km}$ is a typical SMNS equatorial radius."202 The surface magnetic field B.=LOPEBoys is uncoustrained from imo0del calculations aud can be considered a free parameter.," The surface magnetic field $B_* = 10^{13} B_{*,13} {\rm G}$ is unconstrained from model calculations and can be considered a free parameter."203" Equivaleutlv. we may take fas the free parameter. and vary £,=E,/f, accordingly with E, fixed as in eq.1: the fiducial value we choose below ist,~L20davs Guiplviug L,~2.5s10/06s. 1) so as to be cousistent with observations of GRDB991216 (see rofseciafteuv))."," Equivalently, we may take $t_p$ as the free parameter, and vary $L_p=E_p/t_p$ accordingly with $E_p$ fixed as in \ref{eqn:ep}; the fiducial value we choose below is $t_p \sim 120 {\rm days}$ (implying $L_p \sim 2.3 \times 10^{46} {\rm erg \ s^{-1}}$ ) so as to be consistent with observations of GRB991216 (see \\ref{sec:aftenv}) )."204" During fj. L, is expected to be relatively constant. and the wind should enereize a plerionic nebula in the pre-GRB siurouudiugs. a more compact vet much more Iuninous version of the Crab nebula."," During $t_p$, $L_p$ is expected to be relatively constant, and the wind should energize a plerionic nebula in the pre-GRB surroundings, a more compact yet much more luminous version of the Crab nebula."205 The cousequences of such a precursor plerion in the supranova scenario has not been considered previously. and this paper addresses some maportaut dynamical aud radiative effects if may induce. cach providing mrportaut observational diagnostics for the supranova model.," The consequences of such a precursor plerion in the supranova scenario has not been considered previously, and this paper addresses some important dynamical and radiative effects it may induce, each providing important observational diagnostics for the supranova model."206 We discuss the acceleration aud fragmentation of the SN ejecta material by the pleriou-SNR interaction aud its Huplications for Fe line cussion in refsecisnr.. aud inverse Compton scattering of the ambicut pleriou raciation ficld iu the GRB external shock aud the resulting hieh-energy afterglow ciission in refsec:ec..," We discuss the acceleration and fragmentation of the SN ejecta material by the plerion-SNR interaction and its implications for Fe line emission in \\ref{sec:snr}, and inverse Compton scattering of the ambient plerion radiation field in the GRB external shock and the resulting high-energy afterglow emission in \\ref{sec:ec}."207 A brief consideration of the direct detection and identification of the precursor plerion Cluission is given mn rofsecidir.., A brief consideration of the direct detection and identification of the precursor plerion emission is given in \\ref{sec:dir}.208" We will assume a flat lanibda cosmology with Q,,=0.3. Q4=0.7 and I)=τοῖς!1Mpe.+."," We will assume a flat lambda cosmology with $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$ and $H_0=70 {\rm km s^{-1} Mpc^{-1}}$."209 After this paper was subnütted. we became aware of the work of Noumiel and Cranot (2002. hereafter WG02). who also iuvestizated the properties of GRD afterelows occurring inside plerionic nebulae in the contest of the supranova model," After this paper was submitted, we became aware of the work of Könnigl and Granot (2002, hereafter KG02), who also investigated the properties of GRB afterglows occurring inside plerionic nebulae in the context of the supranova model."210 Thev stressed the advantages of such a picture in realizing the relatively ugh maguetic fields aud clectrou injection efficiencies required in the CRB blastwave το explain observed itorvelows. and denmoustrated this by constructing naenetolydrodvuamical (MIID) models for the pleriou under various asstunptious.," They stressed the advantages of such a picture in realizing the relatively high magnetic fields and electron injection efficiencies required in the GRB blastwave to explain observed afterglows, and demonstrated this by constructing magnetohydrodynamical (MHD) models for the plerion under various assumptions."211" Tere we are interested iu some characteristic observable effects caused by the precursor enonie activity wlich are peculiar to the supranova uodel aud through which the model iav be tested by ""ture observations.", Here we are interested in some characteristic observable effects caused by the precursor plerionic activity which are peculiar to the supranova model and through which the model may be tested by future observations.212 The enmiphasis of this paper will be ou the relevant radiative processes aud their observational iuplications: we choose to keep the discussion of the Xenon dynamics relatively simple aud leave more detailed nodeling of this aspect to future studies., The emphasis of this paper will be on the relevant radiative processes and their observational implications; we choose to keep the discussion of the plerion dynamics relatively simple and leave more detailed modeling of this aspect to future studies.213 A study following our approach but with a more realistic treatiuneut of the λαο has receutly heen carried out by Coretta Cranot (2002)., A study following our approach but with a more realistic treatment of the plerion has recently been carried out by Guetta Granot (2002).214 Ax excuplary case of the cdvuamical interaction between a plerion aud a SNR cau be secu in the well-studied Crab uebula. which is known to be accelerating and fragimoeutiug the surrounding SNR. resulting iu the prominent optical fibuneuts instead of a clear shell (e.g. Davidson Fesen 1985. Tester et al.," An exemplary case of the dynamical interaction between a plerion and a SNR can be seen in the well-studied Crab nebula, which is known to be accelerating and fragmenting the surrounding SNR, resulting in the prominent optical filaments instead of a clear shell (e.g. Davidson Fesen 1985, Hester et al."215 1996)., 1996).216 Iu this respect. the powerful supranova pleriou can be even more effective than the Crab.," In this respect, the powerful supranova plerion can be even more effective than the Crab."217" We model the pleriou in a simple wav followine Pacini Salvati (1973. hereatter PS73: see also Bauclicra. DPacni Salvati 1981. Chevalier Revuolds 1981. Chevalier 2000). considering a homogeneous. spherical bubble mto which euergv is injected at a constant rate L,10/9L,,orgstoa fraction £p=Q0.5£pas going into magnetic field aud the rest &=1£p—0.56.95 iuto relativistic electrons."," We model the plerion in a simple way following Pacini Salvati (1973, hereafter PS73; see also Bandiera, Pacini Salvati 1984, Chevalier Reynolds 1984, Chevalier 2000), considering a homogeneous, spherical bubble into which energy is injected at a constant rate $L_p = 10^{46}L_{p,46} {\rm erg \ s^{-1}}$, a fraction $\xi_B=0.5\xi_{B,0.5}$ going into magnetic field and the rest $\xi_e=1-\xi_B=0.5\xi_{e,0.5}$ into relativistic electrons."218 The electrous here are mostly radiative (see refseciec)). and we neglect their pressure for simplicity.," The electrons here are mostly radiative (see \\ref{sec:ec}) ), and we neglect their pressure for simplicity."219" As the plerionic bubble initially plows through the expanding core aud envelope of the progenitor star. it should accelerate the swept-up ejecta material (Chevalier 1977. Chevalier Fransson 1992. Ίνα),"," As the plerionic bubble initially plows through the expanding core and envelope of the progenitor star, it should accelerate the swept-up ejecta material (Chevalier 1977, Chevalier Fransson 1992, KG02)."220" If a fair fraction of the total pleriou cucrey E,~LOere cau be conveved to SN ejecta of several AL... their final attained velocities could reach ον~0.05Ole. cousistent with those imferred fron the observed width of the Fe line in CRD991216 (Piro et al."," If a fair fraction of the total plerion energy $E_p \simeq 10^{53} {\rm erg}$ can be conveyed to SN ejecta of several ${\rm M_\odot}$, their final attained velocities could reach $v_s \sim 0.05-0.1c$, consistent with those inferred from the observed width of the Fe line in GRB991216 (Piro et al."221 2000)., 2000).222 As in the Crab. the SN. ejecta should not eutirelv remain as a spherical shell during this acceleration phase due to Ravleigh-Tavlor (RT) instabilities operating at the pleriou-SNR interface.," As in the Crab, the SN ejecta should not entirely remain as a spherical shell during this acceleration phase due to Rayleigh-Taylor (RT) instabilities operating at the plerion-SNR interface."223" The growth timescale of the RT instability on a spatial scale Rois ter~(RRy. aud since: Ro~;dxfCpp/M.2 pe=DDoreὃπ being: the plerion: magnetic pressure. its ratio to the expansion timescale f,apovRUs Is (Bandiera. Pacini Salvati 1983). where AM. c; and LE.-—Mie2--LOE,org are the SNR ejecta lass. velocity and kinetic energv. aud fig is time after the SN in units of 10 days."," The growth timescale of the RT instability on a spatial scale $R$ is $t_{RT} \sim (R/\ddot R)^{1/2}$, and since $\ddot R \sim 4\pi R^2 p_B/M_s$, $p_B = B_p^2/8\pi$ being the plerion magnetic pressure, its ratio to the expansion timescale $t_{exp} \sim R/v_s$ is (Bandiera, Pacini Salvati 1983), where $M_s$, $v_s$ and $E_s = M_s v_s^2/2 = 10^{51} E_{s,51} {\rm erg}$ are the SNR ejecta mass, velocity and kinetic energy, and $t_{10}$ is time after the SN in units of 10 days."224 Iu the case of the Crab. ter/terpyzl and the RT instability is currently setting iu. consistent with other lines of evidence (Tester et al.," In the case of the Crab, $t_{RT} / t_{exp} \gtrsim 1$ and the RT instability is currently setting in, consistent with other lines of evidence (Hester et al."225 1996. Sankt Tester 1997).," 1996, Sankrit Hester 1997)."226 IToxcever. it can potentially develop much nore rapidly for the supranova plerion.," However, it can potentially develop much more rapidly for the supranova plerion."227 The final outcome of the RT instability is not easv ο. predict. as it will depend on the instability’ non- behavior asx well as other processes such as radiative cooling. and thermal couduction.," The final outcome of the RT instability is not easy to predict, as it will depend on the instability's non-linear behavior as well as other processes such as radiative cooling and thermal conduction."228" We refer to he work of Jun (1998) as a guideline. who carried out 2-dimensioual numerical siuulatious of the pleriou-SNR interaction aud associated RT instabilities. choosing Crab- paraicters for £, aud f aud ucelecting cooling."," We refer to the work of Jun (1998) as a guideline, who carried out 2-dimensional numerical simulations of the plerion-SNR interaction and associated RT instabilities, choosing Crab-like parameters for $L_p$ and $t$ and neglecting cooling."229 Ilis results demoustrate that the RT iustabilitv leads to dramatic effects at late times. strongly disuptiue the SNR shell aud transforming it into pronounced RT “fingers” which carry the majority of the SN ejecta mass aud kinetic," His results demonstrate that the RT instability leads to dramatic effects at late times, strongly disrupting the SNR shell and transforming it into pronounced RT “fingers” which carry the majority of the SN ejecta mass and kinetic"230proximity of a source to the nearest cusp of the tangential caustic.,proximity of a source to the nearest cusp of the tangential caustic.231 This opening angle A@ is measured between Lines joining the centre of the lens to the two outer images οἱ and C., This opening angle $\Delta\theta$ is measured between lines joining the centre of the lens to the two outer images $A$ and $C$.232" As the source moves outwards (towards the nearest cusp). AG»OL fosa7Xe and Raw, will go to zero asymptotically."," As the source moves outwards (towards the nearest cusp), $\Delta\theta \rightarrow 0$, $\mu_{\rm total} \rightarrow233\infty$, and $\Rcusp$ will go to zero asymptotically."234 This relationship holds for any smooth lens potential., This relationship holds for any smooth lens potential.235 Fig., Fig.236 2 illustrates how perturbing structures change the cusp-caustic relation., \ref{fig:Sub_Rcusp_ilus} illustrates how perturbing structures change the cusp-caustic relation.237" The upper panels show the critical curves in the image plane. and the bottom panels are contour maps of Rew), for sources within the tangential caustic in the source plane."," The upper panels show the critical curves in the image plane, and the bottom panels are contour maps of $\Rcusp$ for sources within the tangential caustic in the source plane."238 Left and right. columns show smooth lens potentials and lens potentials with substructures. respectively.," Left and right columns show smooth lens potentials and lens potentials with substructures, respectively."239 Substructures located. near the critical curve will allect images nearby and result. in significantly larger values of Reus. violating the predicted. ratios of image magnifications (Iluxes) given by Iq. (3)).," Substructures located near the critical curve will affect images nearby and result in significantly larger values of $\Rcusp$, violating the predicted ratios of image magnifications (fluxes) given by Eq. \ref{eq:Rcusp}) )."240 Alultiple images of lensecl quasars with small A@ are ideal cases (oO examine violations of the cusp-caustic relation and can be used to put constraints on the properties. of perturbing structures., Multiple images of lensed quasars with small $\Delta\theta$ are ideal cases to examine violations of the cusp-caustic relation and can be used to put constraints on the properties of perturbing structures.241 This is especially true when their Huxes are measured in the radio ancl mid-infrared. as the interpretation of optical and near-infrared Dux ratios is complicated by stellar microlensing and dust extinction.," This is especially true when their fluxes are measured in the radio and mid-infrared, as the interpretation of optical and near-infrared flux ratios is complicated by stellar microlensing and dust extinction."242 At the present time. only five cusp-geometry. lensing systems with image opening angle A@<90° are known.," At the present time, only five cusp-geometry lensing systems with image opening angle $\Delta\theta \leqslant 90^{\circ}$ are known."243 These were used for statistical comparisons to the simulations in our previous work (??7)).," These were used for statistical comparisons to the simulations in our previous work \citealt{Dandan09AquI,244 Dandan2010AqII}) )."245" ALL five cases have surprisingly large Row, values which are dillieult to explain with simple smooth lens models.", All five cases have surprisingly large $\Rcusp$ values which are difficult to explain with simple smooth lens models.246 Of these five (lux-ratio measurements). two that were obtained in the optical have been proven to be allectecl by microlensing: the other three were from the CLASS survey (??)) at radio wavelengths and are thought to be more secure cases of perturbations due to substructures in the lens.," Of these five (flux-ratio measurements), two that were obtained in the optical have been proven to be affected by microlensing; the other three were from the CLASS survey \citealt{Browne2003,Myers2003}) ) at radio wavelengths and are thought to be more secure cases of perturbations due to substructures in the lens."247 Table 2 o£ ? lists all of the currently. observed Ronap- A@ pairs for systems with four distinct point-like images of quasars lensed by one single galaxy., Table 2 of \citet{Chen2011CuspViolation} lists all of the currently observed $\Rcusp$ $\Delta\theta$ pairs for systems with four distinct point-like images of quasars lensed by one single galaxy.248 Our Table 1 lists those with their flux ratios measured in the radio and image opening angles A@x120°.," Our Table \ref{tab:obs120} lists those with their flux ratios measured in the radio and image opening angles $\Delta\theta \leqslant249120^{\circ}$."250 We compare our simulations with this observational sample of lenses. with no additional selection criteria.," We compare our simulations with this observational sample of lenses, with no additional selection criteria."251 Notice that. four of the five lenses listed. in. Table 1l. namely. D0712]472. D1422|231. DB2045]|265 and ALGO4141053. are reported: as having visible companions (satellites/eroup galaxies) projected near the main lensing galaxios ο.," Notice that four of the five lenses listed in Table 1, namely, B0712+472, B1422+231, B2045+265 and MG0414+053, are reported as having visible companions (satellites/group galaxies) projected near the main lensing galaxies \citealt{Fassnacht02B0712Group,252 HB94B1422,Grant04B1422Group,McKean2007,Falco97MG0414}) )."253 Fitting the observed image positions using a singular isothermal ellipsoidal model vields velocity dispersions e ranging from 200+400 kni/s. ancl axis ratios q between 0.7509 (2)). except for B1555. which requires σξ133 km/s and q=0.45 (?)).," Fitting the observed image positions using a singular isothermal ellipsoidal model yields velocity dispersions $\sigma$ ranging from $200-400$ km/s, and axis ratios $q$ between $0.75-0.9$ \citealt{Sluse2012COSMOGRAIL25}) ), except for B1555, which requires $\sigma=133$ km/s and $q=0.45$ \citealt{Marlow1999B1555}) )."254 Given a simulated.lensing system. we compare to the observations in Table 1 by generating a large number of realisations of background. sources. with A@<1207.," Given a simulatedlensing system, we compare to the observations in Table 1 by generating a large number of realisations of background sources with $\Delta\theta \leqslant 120^{\circ}$."255" We calculate. Row, for cach realisation anc evaluate ZRowp{APx2.57) for this ensemble of realisations."," We calculate $\Rcusp$ for each realisation and evaluate $P(\geqslant \Rcusp |256\Delta\theta \pm2.5^{\circ})$ for this ensemble of realisations."257 Ες is defined as the probability for Zeus. measured. for sources with image opening angles CA8257.A0|25°] (ie. within a [ive-degree opening-angle— span centred. at AA) to be larger than a particular threshold. value.," This is defined as the probability for $\Rcusp$, measured for sources with image opening angles $\in258[\Delta\theta-2.5^{\circ},~\Delta\theta+2.5^{\circ}]$ (i.e. within a five-degree opening-angle span centred at $\Delta\theta$ ) to be larger than a particular threshold value."259" Lenses with more perturbations will result in large £u, values for many source positions and thus have a higher PCtas932.57) than lenses with fewer perturbations.", Lenses with more perturbations will result in large $\Rcusp$ values for many source positions and thus have a higher $P(\geqslant \Rcusp | \Delta\theta \pm2.5^{\circ})$ than lenses with fewer perturbations.260 We illustrate our use of the P(esi+2.57) in lig. 3..," We illustrate our use of the $P(\geqslant \Rcusp | \Delta\theta261\pm2.5^{\circ})$ in Fig. \ref{fig:SIEandAq}."262 The top panel shows a typical example of a close triple image configuration for cusp sources with A@x1207., The top panel shows a typical example of a close triple image configuration for cusp sources with $\Delta\theta \leqslant 120^{\circ}$ .263" In this case. the lensing galaxy has a (smooth) singular isothermal cllipsoidal (SUE) profile (sce ο for notations of bj. bse. qa. ancl sy herebelow) with lensing strength b;=0.6"" and axis ratio d4=0.8. and is located at redshift cy=0.6: the source redshift is ος=2."," In this case, the lensing galaxy has a (smooth) singular isothermal ellipsoidal (SIE) profile (see \citet{KKprofile1998} for notations of $b_I$, $b_{\rm SIE}$, $q_3$, and $s_0$ herebelow) with lensing strength $b_I=0.6\arcsec$ and axis ratio $q_3=0.8$, and is located at redshift $z_d=0.6$; the source redshift is $z_s=2$."264 The corresponding contour mapof οίHos[A6£2.57) in the Bau plane is given in the middle panel.," The corresponding contour mapof $P(\geqslant265\Rcusp | \Delta\theta \pm2.5^{\circ})$ in the $\Rcusp$ $\Delta\theta$ plane is given in the middle panel."266 Also plotted: are the radio measurements for the currently best available sample (listed in Table 1))., Also plotted are the radio measurements for the currently best available sample (listed in Table \ref{tab:obs120}) ).267" These are clearly inconsistent with the smooth-lens ow, distribution.", These are clearly inconsistent with the smooth-lens $\Rcusp$ distribution.268 When we include the substructures within the lensing galaxy and its dark matter halo. the regular Re) distribution for a smooth lens potential disappears.," When we include the substructures within the lensing galaxy and its dark matter halo, the regular $\Rcusp$ distribution for a smooth lens potential disappears."269 The bottom panel in Fig., The bottom panel in Fig.270 3. shows the average distribution of PGRow[A632.5) when including the subhalo population from the Aquarius simulations (?2)).," \ref{fig:SIEandAq}271 shows the average distribution of $P(\geqslant \Rcusp | \Delta\theta272\pm2.5^{\circ})$ when including the subhalo population from the Aquarius simulations \citealt{Dandan09AquI}) )."273 At small 0. violations are more significant than on larger scales.," At small $\Delta\theta$ , violations are more significant than on larger scales."274" Ehe smallest Reis, nmieasured among all observed: euspecaustic systems is 0.187 (from Bl422).", The smallest $\Rcusp$ measured among all observed cusp-caustic systems is 0.187 (from B1422).275 In 2.. we caleulated PP(7cnp) which is the probability for. Row to be larger than or equal to Q.18T. computed over all realizations with A@x:907.," In \citet{Dandan09AquI}, we calculated $P^{90}(\Rcusp^{0.187})$, which is the probability for $\Rcusp$ to be larger than or equal to 0.187, computed over all realizations with $\Delta\theta \leqslant 90^{\circ}$."276 PORTE was found to be ~10%.., $P^{90}(\Rcusp^{0.187})$ was found to be $\sim$.277 We concluded that it is cillicult to explain the observed Aes distribution (especially at larger A) with a subhalo population similar to that. produced in the Aquarius simulations., We concluded that it is difficult to explain the observed $\Rcusp$ distribution (especially at larger $\Delta\theta$ ) with a subhalo population similar to that produced in the Aquarius simulations.278 This motivates the search for other sources of perturbations to the lens potential., This motivates the search for other sources of perturbations to the lens potential.279" In this. work. we use P""!:(?chop) as an overall estimate for the probability of observing cusp-caustic violations. in order to compare with our previous work."," In this work, we use $P^{90}(\Rcusp^{0.187})$ as an overall estimate for the probability of observing cusp-caustic violations, in order to compare with our previous work."280 Ehe value of 90° is chosen in orderto compare with ?.. who adopted A6s90 to select. cusp-like lenses which canbe best used to test the cusp-caustic violation.," The value of $90^{\circ}$ is chosen in orderto compare with \citet{AB06mn}, , who adopted $\Delta\theta281\leqslant 90^{\circ}$ to select cusp-like lenses which canbe best used to test the cusp-caustic violation."282 Varving the upper limit of A@ will change the probability of cases with Rep2 0.187., Varying the upper limit of $\Delta\theta$ will change the probability of cases with $\Rcusp\geqslant0.187$ .283 However.," However,"284we can describe the nuuber hnuuinositv per energv of escaping electrous and positrons as where Or) is the stepfunction.,we can describe the number luminosity per energy of escaping electrons and positrons as where $\Theta (x)$ is the step.285 There is another fiux component which should be taken into account., There is another flux component which should be taken into account.286 The electrons/positrons which have the ΟΠΟΙΟΥ lower than τρ(0) ave confined in the SNR and lose their energy adiabatically (Ptuskin Zirakaslivili 2005)., The electrons/positrons which have the energy lower than $\varepsilon_{\rm esc}(t)$ are confined in the SNR and lose their energy adiabatically (Ptuskin Zirakashvili 2005).287 As za) decreases with time. some of the coufined and adiabatically cooled electrons/positrous can escape the shock surface when their energv becomes greater than zí0).," As $\varepsilon_{\rm esc}(t)$ decreases with time, some of the confined and adiabatically cooled electrons/positrons can escape the shock surface when their energy becomes greater than $\varepsilon_{\rm esc}(t)$."288 First. the CR electron/positron nunuber per cherey confined iu the SNR cau be written as where z is the energy of the electrons/positrous at the time #. and it is οι at the time f.," First, the CR electron/positron number per energy confined in the SNR can be written as where $\varepsilon_e^{\prime}$ is the energy of the electrons/positrons at the time $t^{\prime}$, and it is $\varepsilon_e$ at the time $t$."289 The adiabatic loss is determined by the expansion law of the SNR. where BHuxg aud Rex are the radius of the SNR shell aud its expansion velocity. respectively.," The adiabatic loss is determined by the expansion law of the SNR, where $R_{\rm SNR}$ and $\dot{R}_{\rm SNR}$ are the radius of the SNR shell and its expansion velocity, respectively."290" Iu the following we assmue that the SNR is iu the Sedov phase. in which the time dependence of Ryxp is expected to be proportional. to 2f""."," In the following we assume that the SNR is in the Sedov phase, in which the time dependence of $R_{\rm SNR}$ is expected to be proportional to $t^{2/5}$."291" Then the enerey of. electrous/positrous coufined iu the SNR would evolve as aud therefore. the distribution function of confined electrous/positrous cau be evaluated as If the decreasing rate of s, is faster than the adiabatic cooling rate. (Όλοςτ. then a part of confined clectrous/positrous cam escape the SNR shock."," Then the energy of electrons/positrons confined in the SNR would evolve as and therefore, the distribution function of confined electrons/positrons can be evaluated as If the decreasing rate of $\varepsilon_{\rm esc}$ is faster than the adiabatic cooling rate, $-(2/5)\varepsilon_{\rm esc}/t$, then a part of confined electrons/positrons can escape the SNR shock."292 By using the same logic in deriving Eq. (, By using the same logic in deriving Eq. (29321) of Ptuskin Zivakashvili (2005) iu the case of expaucding media. we can evaluate the spectrum of such particles as Ilereafter we consider the spectrum of escapiug clectrous/positrous nuuber luinesity per cuerey as the stun of above two components: We neelect the radiativo enerev loss of clectrous/positrous during the confinement for simplicity.,"21) of Ptuskin Zirakashvili (2005) in the case of expanding media, we can evaluate the spectrum of such particles as Hereafter we consider the spectrum of escaping electrons/positrons number luminosity per energy as the sum of above two components: We neglect the radiative energy loss of electrons/positrons during the confinement for simplicity."294 As we have shown in Eq. (17))," As we have shown in Eq. \ref{escape2}) ),"295" the electron fiux of the second component ON,e2(fe.F) is determined by the difference between the decline rate of the escape energv τς and the euecrgev loss vate of electrous/positrons."," the electron flux of the second component $\dot{N}_{e,{\rm esc},2}(\varepsilon_e,t)$ is determined by the difference between the decline rate of the escape energy $\varepsilon_{\rm esc}$ and the energy loss rate of electrons/positrons."296" Iu the cases shown iu this study. the decline rate of the escape energv ds ~asft3s10ο”... (ax12.6. depending onu the model). while the adiabatic cooling rate aud the radiative cooling rate is ~12&10?(lTeV)(+Lotyy)GeVsecD ~bz?LOMee,1TeVYGeVxocD. respectively. where for the latter we take the cooling rate for the interstellar space (see Sec."," In the cases shown in this study, the decline rate of the escape energy is $\sim \alpha \varepsilon_e/t\sim 3\times 10^{-9}\left( \varepsilon_e/1{\rm TeV}\right) \left(t/10^4{\rm yr}\right) ^{-1}{\rm GeV}~{\rm sec}^{-1}$ $\alpha\simeq 1-2.6$, depending on the model), while the adiabatic cooling rate and the radiative cooling rate is $\sim 1.2\times 10^{-9}\left(\varepsilon_e/1{\rm TeV}\right) \left(t/10^4{\rm yr}\right) ^{-1}{\rm GeV}~{\rm sec}^{-1}$, $\sim b\varepsilon_e^2 \sim 10^{-10}\left(\varepsilon_e/1{\rm TeV}\right) ^2 {\rm GeV}~{\rm sec}^{-1}$, respectively, where for the latter we take the cooling rate for the interstellar space (see Sec."297 2)., 2).298 Therefore. even if we take iuto account the radiative cooling the flux of the clectrous/positrous which have once been confined iu the SNR does not change its order from our calculation.," Therefore, even if we take into account the radiative cooling the flux of the electrons/positrons which have once been confined in the SNR does not change its order from our calculation."299" Moreover. in the interstellar space the diffusion timescale for TeV clectrous can be estimated as and therefore the enerev loss of TeV olectrous during the propagation is at most As,fe.~1.(1ματι) 1054."," Moreover, in the interstellar space the diffusion timescale for TeV electrons can be estimated as and therefore the energy loss of TeV electrons during the propagation is at most $\Delta \varepsilon_e/\varepsilon_e\sim 1-(1+bt_{\rm diff}\varepsilon_e)^{-1}\sim 10\%$ ."300 Then we can sav that oein this enerev range both of the cnerev losses in the PWN aud in the interstellar space are small., Then we can say that in this energy range both of the energy losses in the PWN and in the interstellar space are small.301 However. if the magnetic feld iu the PWN/SNR is strongly amplificcd from the interstellar value. the cenerev loss rate due to the svuchrotrou Cluission nav be faster than that due to the adiabatic expansion of the SNR. aud even than that due to the dechne rate of the escape energv.," However, if the magnetic field in the PWN/SNR is strongly amplified from the interstellar value, the energy loss rate due to the synchrotron emission may be faster than that due to the adiabatic expansion of the SNR, and even than that due to the decline rate of the escape energy."302 Iu such case. the confined electrous/positrous cannot escape the SNR later and ouly the first component Noe.) would be emitted and observed at the Earth.," In such case, the confined electrons/positrons cannot escape the SNR later and only the first component $\dot{N}_{e,{\rm esc},1}$ would be emitted and observed at the Earth."303" Αννας, as either the streneth of the magnetic field iu the PWN/SNR nor its tine evolution is ecucrally uncertain. it is difficult"," Anyway, as either the strength of the magnetic field in the PWN/SNR nor its time evolution is generally uncertain, it is difficult"304our model may be far [rom sufficient for a well-correlated distribution.,our model may be far from sufficient for a well-correlated distribution.305 Our assumption is reinforce by the [act that forecasts for [ree atinosphere seeing are generally coucentrated in a uarrow region despite their mean is close to that of observation. as shown in Figure 6..," Our assumption is reinforce by the fact that forecasts for free atmosphere seeing are generally concentrated in a narrow region despite their mean is close to that of observation, as shown in Figure \ref{fig-6}."306 This phenomenou implies that turbulence over sub-kilometer scale in the vertical direction uuelt be the major contributor toa bad seeing condition iu [ree atmosphere region., This phenomenon implies that turbulence over sub-kilometer scale in the vertical direction might be the major contributor to a bad seeing condition in free atmosphere region.307 Tu final. we compare our result with several major models/forecasters (Table 8)).," In final, we compare our result with several major models/forecasters (Table \ref{tbl-7}) )."308" The RMSE and «30%.) error probabilities⋅⋅⋅ for⋅ the original⋅⋅ ANPr model. Verniu-TatarskiM. model. Cs,4? /seeiug4. ujeciau aid seeing mean are derived by Trinquet Vernin with their experimental pm‘ofile observations."," The RMSE and $<30$ error probabilities for the original AXP model, Vernin-Tatarski model, $C_{N}^{2}$ /seeing median and seeing mean are derived by Trinquet Vernin with their experimental profile observations."309 The MAWC/WRE ιούς) is initiated with the GES model output. but the simulation would eveutually derives output with final grid spacing of Im aud a vertical laver number of 10. resulting a forecast RMSE at 0.367 for the Ist night.," The MKWC/WRF model is initiated with the GFS model output, but the simulation would eventually derives output with final grid spacing of 1km and a vertical layer number of 40, resulting a forecast RMSE at 0.36"" for the 1st night."310 As revealed by Table 8.. the RMSE uncertaiuty ranges of GES/ANP model just fall between the mean of MAWC forecaster aud the original AXP model. while the <30% error probability is in the better cluster in all moclels wuder most. cases.," As revealed by Table \ref{tbl-7}, the RMSE uncertainty ranges of GFS/AXP model just fall between the mean of MKWC forecaster and the original AXP model, while the $<30$ error probability is in the better cluster in all models under most cases."311 Interestingly. a “direct” comparison of the CES/ANP model with the MINXWC' forecaster over the seeiug forecast for the same site (Mauna Wea) even slightly favors the former (0.267. versus 0.287 for three-night mean).," Interestingly, a “direct” comparison of the GFS/AXP model with the MKWC forecaster over the seeing forecast for the same site (Mauna Kea) even slightly favors the former (0.26"" versus 0.28"" for three-night mean)."312 Iu short. this result has highlighted the potential of the GES/ANP moclel to be a competitive forecast tool once the laver degeneracy aud high-altitude issues are solved.," In short, this result has highlighted the potential of the GFS/AXP model to be a competitive forecast tool once the layer degeneracy and high-altitude issues are solved."313 We have carried out a coimpreliensive study over the topic of performing automatic nuimeric forecast of cloud cover ancl atmospheric seeing with the Global Forecast System. au operational elobal model.," We have carried out a comprehensive study over the topic of performing automatic numeric forecast of cloud cover and atmospheric seeing with the Global Forecast System, an operational global model."314 Sequence observations on cloud. cover aud. atmospheric seeing from 9 sites from different regious of the world with different climatic background iu the period of January 2008 to December 2009 are used to evaluate the forecast., Sequence observations on cloud cover and atmospheric seeing from 9 sites from different regions of the world with different climatic background in the period of January 2008 to December 2009 are used to evaluate the forecast.315 Although the performance of the model forecast may not be comparable with the humau-participated forecast. our study has shown that the forecast to be acceptable lor basic observing relerence.," Although the performance of the model forecast may not be comparable with the human-participated forecast, our study has shown that the forecast to be acceptable for basic observing reference."316 Our study has also reveal tlie possibility to gain better performance from the moclel with additional ellorts on model refinement., Our study has also reveal the possibility to gain better performance from the model with additional efforts on model refinement.317 For cloud cover forecast. we have fouud that the proportion of perfect [orecast varies from 50% to cδὲ or all three sites we evaluated. including a site located in subtropical region with a very humid climate (Lulin).," For cloud cover forecast, we have found that the proportion of perfect forecast varies from $\sim50$ to $\sim85$ for all three sites we evaluated, including a site located in subtropical region with a very humid climate (Lulin)."318 Iu particular. we have found that the moclel is capable to detect a significant amount of occurring clouds while the false alarm rate is moderate.," In particular, we have found that the model is capable to detect a significant amount of occurring clouds while the false alarm rate is moderate."319 The probability of detection is still measured to be even for site with very low cloudy probability (Paranal)., The probability of detection is still measured to be even for site with very low cloudy probability (Paranal).320 For atmospheric seeing forecast. we adopted the ANP model introduced by Tringuet Vernin.," For atmospheric seeing forecast, we adopted the AXP model introduced by Trinquet Vernin."321 We found that forecast for eutire atinosphere tends to slightly overestimate the seeing. while the [ree atinosphere forecast teuds to the opposition.," We found that forecast for entire atmosphere tends to slightly overestimate the seeing, while the free atmosphere forecast tends to the opposition."322 The RMSE for [ree atinosphere seeing is smaller (0.22°-0.12°) than that of entire atmosphere (0.26°-0.507). but both values cau indicate a decent quality of the forecast comparing with the other major models.," The RMSE for free atmosphere seeing is smaller (0.22""-0.42"") than that of entire atmosphere (0.26""-0.50""), but both values can indicate a decent quality of the forecast comparing with the other major models."323 Further aualysis suggests that a, Further analysis suggests that a324speed of the interstellar eas.,speed of the interstellar gas.325 These assimuptionus lead to secular decrease of semianajor axis e of the particle., These assumptions lead to secular decrease of semi-major axis $a$ of the particle.326 The secular time derivative of the semi-auajor axis is negative aud proportional to , The secular time derivative of the semi-major axis is negative and proportional to $a$.327This result is not iu accordance with Scherer (2000) who has stated that the semi-major of the particle incrases exponeutiallv., This result is not in accordance with Scherer (2000) who has stated that the semi-major of the particle increases exponentially.328 Scherers statement is Incorrect and our analytical result is confiiied by our detailed nmnuerica inteeration of equation of motion (see also Fig., Scherer's statement is incorrect and our analytical result is confirmed by our detailed numerical integration of equation of motion (see also Fig.329 1)., 4).330" If the hwdrogeu gas velocity vector wy lies iu the article's orbital plane aud the major axis of the orbit i not perpendicular to 9j. then the product of (secular) eccentricity and iaguitude of the radial component of e; neasured in peribelion is. approximately, constant during orbital evolution."," If the hydrogen gas velocity vector $\vec{v}_{H}$ lies in the particle's orbital plane and the major axis of the orbit is not perpendicular to $\vec{v}_{H}$, then the product of (secular) eccentricity and magnitude of the radial component of $\vec{v}_{H}$ measured in perihelion is, approximately, constant during orbital evolution."331 We considered simmltaueous action of the P-R effect. he radial solar wind aud the interstellar eas flow.," We considered simultaneous action of the P-R effect, the radial solar wind and the interstellar gas flow."332 Niunerical iucerations showed that the action of the flow of interstellar eas cau be more important than the action of the electromagnetic and corpuscular raciation of the Sun. as for the motion of dust particles orbiting the Siu iu outer warts of the Solar System (see Fig.," Numerical integrations showed that the action of the flow of interstellar gas can be more important than the action of the electromagnetic and corpuscular radiation of the Sun, as for the motion of dust particles orbiting the Sun in outer parts of the Solar System (see Fig."333 5)., 5).334 Plysical decrease of senianajor axis can be more than 2-tines ereater than the value produced bv the P-R effect and radial solar wind., Physical decrease of semi-major axis can be more than 2-times greater than the value produced by the P-R effect and radial solar wind.335 The evolution of ecceutriicitv can also be an increasing functiou of time when we cousider the P-R effect aud the radial solar wiud together with the flow of the neutral interstellar eus., The evolution of eccentricity can also be an increasing function of time when we consider the P-R effect and the radial solar wind together with the flow of the neutral interstellar gas.336ons This is also relevaut difference from the action of the P-R effect aud the radial solar wind when secular decrease of eeceutricity occurs., This is also relevant difference from the action of the P-R effect and the radial solar wind when secular decrease of eccentricity occurs.337 Simultaneous action of all τος effects vields that the secular time derivative of the argument of perihelion may not he equal to zero. in general.," Simultaneous action of all three effects yields that the secular time derivative of the argument of perihelion may not be equal to zero, in general."338 Gravitation of four major planets was also directly added into the equation of motion. see Eq. (," Gravitation of four major planets was also directly added into the equation of motion, see Eq. ("33915).,48).340 This access correctly describes capture of dust evais into nie motion resonances with the planets., This access correctly describes capture of dust grains into mean motion resonances with the planets.341 Our physical approach differs from the Scherer’s approach (Scherer 20003. who has used some kind of secular access to eravitational influence of the planets.," Our physical approach differs from the Scherer's approach (Scherer 2000), who has used some kind of secular access to gravitational influence of the planets."342 Asstuuption on an existence of dust rine in the zone of the Edgeworth-INuiper bel is du contradiction with rapid iucrease of ecceutricity of the ving due to an acceleration caused|w the interstellar eas flow., Assumption on an existence of dust ring in the zone of the Edgeworth-Kuiper belt is in contradiction with rapid increase of eccentricity of the ring due to an acceleration caused by the interstellar gas flow.343 Speed of he eeceutricity mnerease (time derivative of eccentricity) is roughly inversely proportional to the particles size aud nass density., Speed of the eccentricity increase (time derivative of eccentricity) is roughly inversely proportional to the particle's size and mass density.344 As the eccentricity of the particles increases. he particles approach the planets.," As the eccentricity of the particles increases, the particles approach the planets."345 The particles in the ring are under the eravitational iuflueuce of the plaucts., The particles in the ring are under the gravitational influence of the planets.346 The particles evolve alsoin semiauajor axis aud they cau jo temporarily captured iuto a mean motion resonance., The particles evolve also in semi-major axis and they can be temporarily captured into a mean motion resonance.347 The particles cau remain im chaotic orbits between orbits of the planets. or. the particles are ejected to high eccentric orbits due to close encounters with oue of the plaucts.," The particles can remain in chaotic orbits between orbits of the planets, or, the particles are ejected to high eccentric orbits due to close encounters with one of the planets."348 Ouly particles with greater size and mass deusity should survive in the dust ring for a long time., Only particles with greater size and mass density should survive in the dust ring for a long time.349 A relevant result of the paper is that equation of motion in the form of Eq. (, A relevant result of the paper is that equation of motion in the form of Eq. (350LL) Cor. Eq.,"44) (or, Eq."351 I8) and Ec4 (, 48) and Eqs. (352L5}-€11) have to be used im modehue of orbital evolution of dust grains in the Solar Systeu.,45)-(46) have to be used in modeling of orbital evolution of dust grains in the Solar System.353 The infinence of the fast interstellar ucutra eas flow night not be ignored iu general investigations on evolution of dust particles in the zone of the Edgeworth-Ixuiper belt., The influence of the fast interstellar neutral gas flow might not be ignored in general investigations on evolution of dust particles in the zone of the Edgeworth-Kuiper belt.354(at a distance of about 2 pe).,(at a distance of about 2 pc).355 Given its spectral characteristics. which will be shown later in the paper. s4 might also be an evolved star: a galactic AGB or an AGB in M33 at a large galactocentric distance. beyond the area surveyed by ?..," Given its spectral characteristics, which will be shown later in the paper, s4 might also be an evolved star: a galactic AGB or an AGB in M33 at a large galactocentric distance, beyond the area surveyed by \citet{2007ApJ...664..850M}."356 We keep the evolved stars in our sample to check if there is any detectable CO emission from the surrounding region., We keep the evolved stars in our sample to check if there is any detectable CO emission from the surrounding region.357 Some CO Jz1-0 emission has already been detected with a 45 aresee wide beam (FCRAO) around a few sources of our sample (?).., Some CO J=1-0 emission has already been detected with a 45 arcsec wide beam (FCRAO) around a few sources of our sample \citep{2004ApJ...602..723H}.358 Given the large beam size. we do not know if the detected gas is associated to or in the proximity of the sources.," Given the large beam size, we do not know if the detected gas is associated to or in the proximity of the sources."359 Therefore. we have searched for CO emission using the smaller IRAM-30 mt beam from all the sources in Table 1.," Therefore, we have searched for CO emission using the smaller IRAM-30 mt beam from all the sources in Table 1."360 The CO J=1-0 and J=2-1! lines have been observed during August 2007 with a FWHM beam of 24 aresee at 115 GHz and of 12 aresee at230 GHz., The CO J=1-0 and J=2-1 lines have been observed during August 2007 with a FWHM beam of 24 arcsec at 115 GHz and of 12 arcsec at 230 GHz.361 At 24 all the sources are smaller in size than the telescope beam at 230 GHz., At 24 all the sources are smaller in size than the telescope beam at 230 GHz.362 We have observed the sources in position switching mode. using the receiver combination AT00/BIOO and A230/B230 and the VESPA backend system with 240 MHz bandwidth (320 kHz resolution).," We have observed the sources in position switching mode, using the receiver combination A100/B100 and A230/B230 and the VESPA backend system with 240 MHz bandwidth (320 kHz resolution)."363 One source was centered on the ON position. another source in the OFF position.," One source was centered on the ON position, another source in the OFF position."364 The OFF source is chosen in a region with different line of sight velocity than the ON position (as seen through 21-em maps)., The OFF source is chosen in a region with different line of sight velocity than the ON position (as seen through 21-cm maps).365 The spectra have been smoothed to 1 km s! and the data from both receivers has been averaged., The spectra have been smoothed to 1 km $^{-1}$ and the data from both receivers has been averaged.366 In Table 2 we summarize the CO data: integrated emission I (in units of main beam temperature K km s7! ). mean velocity V. line width W ( full width at half maximum: hereafter. FWHM) and peak intensity P. are estimated using gaussian fits to the lines.," In Table 2 we summarize the CO data: integrated emission I (in units of main beam temperature K km $^{-1}$ ), mean velocity V, line width W ( full width at half maximum; hereafter, FWHM) and peak intensity P, are estimated using gaussian fits to the lines."367 The rms refers to a spectral resolution of 2.2 and 1.1 km s! for the CO J21-0 and J=2-1. respectively.," The rms refers to a spectral resolution of 2.2 and 1.1 km $^{-1}$ for the CO J=1-0 and J=2-1, respectively."368 The line widths have been measured using correlator spectra after correcting for hanning., The line widths have been measured using correlator spectra after correcting for hanning.369 We shall use as CO J=1-0 line intensity the average value of that derived by fitting a gaussian to the detected emission and that obtained by summing the flux in each channel inside the signal window (determined individually for each spectrum)., We shall use as CO J=1-0 line intensity the average value of that derived by fitting a gaussian to the detected emission and that obtained by summing the flux in each channel inside the signal window (determined individually for each spectrum).370 The result is given in Table 3., The result is given in Table 3.371 As uncertainties on the CO line intensities. we consider the largest value between the following ones: uncertainty derived from the gaussian fit. the rms of the spectra integrated over the signal window. the dispersion between the intensity derived from gaussian fit. and the intensity derived by the integral inside the window signal.," As uncertainties on the CO line intensities, we consider the largest value between the following ones: uncertainty derived from the gaussian fit, the rms of the spectra integrated over the signal window, the dispersion between the intensity derived from gaussian fit, and the intensity derived by the integral inside the window signal."372 The gain is 6.3 and 8.7 Jy K! at 110 and 235 GHz. respectively.," The gain is 6.3 and 8.7 Jy $^{-1}$ at 110 and 235 GHz, respectively."373 Using these values. we derive the CO fluxes of the two rotational levels and their ratios.," Using these values, we derive the CO fluxes of the two rotational levels and their ratios."374 In Table 3 we also quote the H» column densities derived. from FCRAO (beam=45 arcsec) integrated CO J=1-0 fluxes. using a H» conversion factor X¢o=2.8 107 emo? ΚΙ.," In Table 3 we also quote the $_2$ column densities derived from FCRAO (beam=45 arcsec) integrated CO J=1-0 fluxes, using a $_2$ conversion factor $_{CO}$ =2.8 $^{20}$ $^{-2}$ $^{-1}$."375 If there are no entries in the column corresponding to the FCRAO column density. it means that the source was off the region mapped by FCRAO.," If there are no entries in the column corresponding to the FCRAO column density, it means that the source was off the region mapped by FCRAO."376 The bolometric luminosity for star forming regions in Table 3 is computed as the sum of the FUV and NUV lummosities uncorrected for absorption. added to the total infrared luminosity (hereafter. TIR luminosity):," The bolometric luminosity for star forming regions in Table 3 is computed as the sum of the FUV and NUV luminosities uncorrected for absorption, added to the total infrared luminosity (hereafter, TIR luminosity):"377cclouds have been observed in shadow against the soft x-ray background Aurrows Alendenhall 1991).,clouds have been observed in shadow against the soft x-ray background Burrows Mendenhall 1991).378 There have even been Claims for x-ray shadows from extragalactic objects (Barber. ltoberts Warwick 1996).," There have even been claims for x-ray shadows from extragalactic objects (Barber, Roberts Warwick 1996)."379 After repeated attempts. we were not allowed access to the ROSAT high resolution database in order to check this.," After repeated attempts, we were not allowed access to the ROSAT high resolution database in order to check this."380 Phe low resolution. (40). ROSAT all-sky survey shows no obvious emission or absorption in the cürection of the Smith Cloud.," The low resolution ), ROSAT all-sky survey shows no obvious emission or absorption in the direction of the Smith Cloud."381 The recent detections of the Magellanic Stream in ((Wetner Williams 1996) provided Bland-Llawthorn Maloney (1998) with the critical normalisation for the emergent UV lux from the Galactic disk., The recent detections of the Magellanic Stream in (Weiner Williams 1996) provided Bland-Hawthorn Maloney (1998) with the critical normalisation for the emergent UV flux from the Galactic disk.382 This model has been used to oediet the ionising field throughout the Galactic halo., This model has been used to predict the ionising field throughout the Galactic halo.383 We determine a distance to the Smith Cloud of 2644 kpe on the xisis of the flux alone., We determine a distance to the Smith Cloud of $\pm$ 4 kpc on the basis of the flux alone.384 In support of this picture. the iis greatly enhanced at the velocity of the cloud and. nunust be very weak. which together indicate a dilute ionising field impinging the cloud surface.," In support of this picture, the is greatly enhanced at the velocity of the cloud and must be very weak, which together indicate a dilute ionising field impinging the cloud surface."385 Possible complications are. cloud. geometry. porosity. and uncertain extinction corrections.," Possible complications are cloud geometry, porosity, and uncertain extinction corrections."386 X lower mean cisk opacity would put the clouds further away., A lower mean disk opacity would put the clouds further away.387 We anticipate that our model is more reliable for clouds at greater vertical distances as this tends to average out structure in the distribution of UV sources., We anticipate that our model is more reliable for clouds at greater vertical distances as this tends to average out structure in the distribution of UV sources.388 Clouds within a few kpe of the Galactic plane could be in relative shadow. particularly. for opaque disk mocels.," Clouds within a few kpc of the Galactic plane could be in relative shadow, particularly for opaque disk models."389 Alore accurate distances for HEVC's will come from large ssurvevs with optical follow-up., More accurate distances for HVCs will come from large surveys with optical follow-up.390 A crucial development has been the availability of target [ists with higher sensitivity and better resolution. in particular. the Parkes AlMulti-beam Survey (Staveley-Smith 1997).," A crucial development has been the availability of target lists with higher sensitivity and better resolution, in particular, the Parkes Multi-beam Survey (Staveley-Smith 1997)."391 More. recent detections at the AAT are sulliciently strong to suggest that a significant number of clouds should be observable with this technique., More recent detections at the AAT are sufficiently strong to suggest that a significant number of clouds should be observable with this technique.392 As such. it provides a crucial test of the Spereel model 11996: 11996) which places roughly of half allhigh velocity clouds at extragalactic distances (1 Alpe) within the Local Group. in which case none of the clouds should be detectable at these levels.," As such, it provides a crucial test of the Blitz-Spergel model 1996; 1996) which places roughly half of all high velocity clouds at extragalactic distances $\sim 1$ Mpc) within the Local Group, in which case none of the clouds should be detectable at these levels."393 We have set out. to. present. a simplified picture as a challenge to theorists and cxperimentalists alike., We have set out to present a simplified picture as a challenge to theorists and experimentalists alike.394 The observational programmes that we have described here only require small to mecitun-sizecl telescopes 4m). since both LIVC€ anc Stream clouds subtenc large. angles. ancl the line surface brightness is expected to be fairly constant over laree angular scales (< 17).," The observational programmes that we have described here only require small to medium-sized telescopes $-$ 4m) since both HVC and Stream clouds subtend large angles, and the line surface brightness is expected to be fairly constant over large angular scales $\leq 1^\circ$ )."395 Since the Fabry-Perot interferogram is binned azimuthally in order to produce the detection. the site does not require good seeing although it should. have relatively good. photometric stability.," Since the Fabry-Perot interferogram is binned azimuthally in order to produce the detection, the site does not require good seeing although it should have relatively good photometric stability."396" A dry site ds favoured because variable water vapour features can complicate sky subtraction for £,« 1 cem"" pe.", A dry site is favoured because variable water vapour features can complicate sky subtraction for $<$ 1 $^{-6}$ pc.397" We encourage a more widespread. interest. in the Fabry-Perot ""staring technique as Hb ds set to make a profound contribution to the understanding of both galactic and oxtragalactic radiation fields over the coming vears.", We encourage a more widespread interest in the Fabry-Perot `staring' technique as it is set to make a profound contribution to the understanding of both galactic and extragalactic radiation fields over the coming years.398 We are indebted. to €. da Costa for elarifving recent developments in the study of dwarf spheroidals., We are indebted to G. da Costa for clarifying recent developments in the study of dwarf spheroidals.399 We thank li. LIbata. D. Ixoribalski and CoG. Pinney for their assistance with important references. and a highly competent. referee [or recommendations which encouraged. us to be more circeumspect in our analysis of errors.," We thank R. Ibata, B. Koribalski and C.G. Tinney for their assistance with important references, and a highly competent referee for recommendations which encouraged us to be more circumspect in our analysis of errors."400since all terms proportional to flexion or g? will be dropped anyhow due to the multiplication with 3.,since all terms proportional to flexion or $g^2$ will be dropped anyhow due to the multiplication with $\bar{\beta}$.401" Then we have: and which lead to and20b,, where the upper index ""t"" has been dropped."," Then we have: and which lead to and, where the upper index ""t"" has been dropped."402 We derive here the relation between ε ellipticity and η ellipticity for a source with elliptical isophotes., We derive here the relation between $\epsilon$ ellipticity and $\eta$ ellipticity for a source with elliptical isophotes.403 Without loss of generality we assume the source profile to be described by an elliptical gaussian with ellipticity e., Without loss of generality we assume the source profile to be described by an elliptical gaussian with ellipticity $\epsilon$.404 Furthermore we select the reference frame such that e?=0. , Furthermore we select the reference frame such that $\epsilon_2 = 0$ 405The atmospheric parameters of 6687 were measured by fitting metal-line blanketed LTE models to the UVES spectra and are given in Lisker et al. (2005)).,The atmospheric parameters of 687 were measured by fitting metal-line blanketed LTE models to the UVES spectra and are given in Lisker et al. \cite{lisker}) ).406 In order to derive v4sin/ and the elemental abundances. we coadded the observed high resolution spectra after shifting them to rest wavelength.," In order to derive $v_{\rm rot}\,\sin{i}$ and the elemental abundances, we coadded the observed high resolution spectra after shifting them to rest wavelength."407 The spectrum was then compared with rotationally broadened. synthetic line profiles calculated with the LINFOR program (developed by Holweger. Steffen and Steenbock at Kiel university. modified by Lemke 1997)).," The spectrum was then compared with rotationally broadened, synthetic line profiles calculated with the LINFOR program (developed by Holweger, Steffen and Steenbock at Kiel university, modified by Lemke \cite{lemke}) )."408 Due to the wide slit used (2.17) for the SPY survey (see Naprwotzki et al., Due to the wide slit used $2.1''$ ) for the SPY survey (see Napiwotzki et al.409 2001 for details) the resolution of the UVES spectra is seeing dependent in most cases., \cite{napiwotzki9} for details) the resolution of the UVES spectra is seeing dependent in most cases.410 In order to measure Vysin; accurately. the instrumental profile has to be taken into account.," In order to measure $v_{\rm rot}\,\sin{i}$ accurately, the instrumental profile has to be taken into account."411 We used the ESO archive to obtain the seeing conditions during the exposure times of the UVES spectra (DIMM seeing monitor. Sarazin Roddier 1990)). took the average (1.13 arcsec) and folded our models with the instrumental profile.," We used the ESO archive to obtain the seeing conditions during the exposure times of the UVES spectra (DIMM seeing monitor, Sarazin Roddier \cite{sarazin}) ), took the average $1.13\,{\rm arcsec}$ ) and folded our models with the instrumental profile."412 The projected rotational velocity was measured simultaneously with the elemental abundances to Vorsing=212x20kms! using the six strongest metal lines.," The projected rotational velocity was measured simultaneously with the elemental abundances to $v_{\rm rot}\sin{i}=21.2\pm2.0\,{\rm km\,s^{-1}}$ using the six strongest metal lines."413 In addition. an abundance and rotational broadening fit was done using two prominent helium lines 4472. 5876) and LTE model spectra (Heber et al. 2000)).," In addition, an abundance and rotational broadening fit was done using two prominent helium lines 4472, 5876) and LTE model spectra (Heber et al. \cite{heber2}) )."414" The resulting v,sin;=19.82.2kms?! turned out to be perfectly consistent with the value measured from the metal lines."," The resulting $v_{\rm rot}\sin{i}=19.8\pm2.2\,{\rm km\,s^{-1}}$ turned out to be perfectly consistent with the value measured from the metal lines."415 Since metal lines are more sensitive to rotational broadening we adopt the value derived from these lines for our analysis., Since metal lines are more sensitive to rotational broadening we adopt the value derived from these lines for our analysis.416 The parameters are given in Table 2.., The parameters are given in Table \ref{tab:par}. .417 Our results depend very much on the accuracy of the Viosinf and logs measurements., Our results depend very much on the accuracy of the $v_{\rm rot}\sin{i}$ and $\log{g}$ measurements.418 A thorough discusstor of the error in surface gravity is given by Lisker et al. (2005))., A thorough discussion of the error in surface gravity is given by Lisker et al. \cite{lisker}) ).419" To quantify the v,sin/ error, we carried out numerical simulations."," To quantify the $v_{\rm rot}\sin{i}$ error, we carried out numerical simulations."420 For this. synthetic spectra with fixed rotational broadening were computed and convolved with the instrumental profile.," For this, synthetic spectra with fixed rotational broadening were computed and convolved with the instrumental profile."421 Random noise was added to mimic the observed spectra., Random noise was added to mimic the observed spectra.422 The rotational broadening was measured i the way described above using a grid of synthetic spectra for various rotational and instrumental broadenings as well as S/ levels., The rotational broadening was measured in the way described above using a grid of synthetic spectra for various rotational and instrumental broadenings as well as S/N levels.423 Variations in the instrumental profile and the noise level were the dominant error sources., Variations in the instrumental profile and the noise level were the dominant error sources.424 Behr (2003)) used a similar method to measure the low vjsin/ of Blue Horizontal Branch stars from high resolution spectra., Behr \cite{behr}) ) used a similar method to measure the low $v_{\rm rot}\sin{i}$ of Blue Horizontal Branch stars from high resolution spectra.425" The errors given im that work are of the same order as the one given here for 6687 (2.0kms7!),"," The errors given in that work are of the same order as the one given here for 687 $2.0\,{\rm km\,s^{-1}}$ )."426 The analysis strategy used ts described only briefly., The analysis strategy used is described only briefly.427 For details we refer the reader to Geter et al. (2007.. 2008)).," For details we refer the reader to Geier et al. \cite{geier1}, \cite{geier2}) )."428" Since the spectrum of 6687 ts single-lined. it contains no informationabout the orbital motion of the companion. and thus only the mass function f,=MÀcompsin?HOMwpMamP=PRAπο "," Since the spectrum of 687 is single-lined, it contains no informationabout the orbital motion of the companion, and thus only the mass function $f_{\rm m} = {M_{\rm comp}^3 \sin^3i}/(M_{\rm comp} + M_{\rm sdB})^2 = 429P K^3/2 \pi G$ "430in this age range.,in this age range.431 Within the time period considered here - JO Myr to 1 Cr - we do not find evidence for cluster dissolution., Within the time period considered here - 10 Myr to 1 Gyr - we do not find evidence for cluster dissolution.432 Iu both galaxies. the clusters become fainter with increasing age.," In both galaxies, the clusters become fainter with increasing age."433 The very massive hot stars. which are still present iu the voune star clusters and contribute most of the light. become famter and redder with mereasiug age and so do the star clusters.," The very massive hot stars, which are still present in the young star clusters and contribute most of the light, become fainter and redder with increasing age and so do the star clusters."434 This trend cau be seen iu both the LMC aud the SAIC., This trend can be seen in both the LMC and the SMC.435 The total cluster luminosity increases with increasing radius due to a larger umber of stars within the cluster radius., The total cluster luminosity increases with increasing radius due to a larger number of stars within the cluster radius.436Although dwarf ealaxies are the inost nunierous extragalactic objects in the nearby Universe (e.e.. Fereusou Bineech 1991: Dingeeli. Saudage Tanuuaun L9sa: Mateo 1998). it seems unclear how they are related in origin to typical large (1.0. £) galaxies.,"Although dwarf galaxies are the most numerous extragalactic objects in the nearby Universe (e.g., Ferguson Binggeli 1994; Binggeli, Sandage Tammann 1988; Mateo 1998), it seems unclear how they are related in origin to typical large (i.e., $L^{*}$ ) galaxies."437 Dwart galaxies could be formed through the same formation mechiauisui as hat of large galaxies: e.g.. gravitational collapse of proogalactic gas clouds (Dekel Silk 1986: White Freuk 1991: Freuk et al.," Dwarf galaxies could be formed through the same formation mechanism as that of large galaxies; e.g., gravitational collapse of protogalactic gas clouds (Dekel Silk 1986; White Frenk 1991; Frenk et al."438 1996: RKauffiiaun. Nusser Steinmetz 1997).," 1996; Kauffmann, Nusser Steinmetz 1997)."439 ITowever. it is known that dsvarf elliptical ealaxies (dEs) apparently belong to a ditfereut class frou normal lavee cllipticals (Es) in the fundamental plane (6.8. Isormendy 1985). sugeesting that the formation and/or evolution processes of dwarfs may not always be the same as those of larger ellipticals.," However, it is known that dwarf elliptical galaxies (dEs) apparently belong to a different class from normal large ellipticals (Es) in the fundamental plane (e.g., Kormendy 1985), suggesting that the formation and/or evolution processes of dwarfs may not always be the same as those of larger ellipticals."440 It las been argued from an observational view poiut that dwiuf galaxies may be formed by galaxy collisions because there appears to be morphological evidence for dwarf galaxies in the tidal tails of interacting galaxies (Zwicky 1956: Scliweizer 1978: Duc et al., It has been argued from an observational view point that dwarf galaxies may be formed by galaxy collisions because there appears to be morphological evidence for dwarf galaxies in the tidal tails of interacting galaxies (Zwicky 1956; Schweizer 1978; Duc et al.441" 2000): Ίνοι, gas-vich dwarf irregular galaxies. can be made out of stellar and gaseous material pulled out into intergalactic space by tidal forces from the disks of colliding pareut galaxies."," 2000); i.e., gas-rich dwarf irregular galaxies, can be made out of stellar and gaseous material pulled out into intergalactic space by tidal forces from the disks of colliding parent galaxies."442 This possibility has been receutlv reinforced by a number of pieces of observational evidence (Sclaveizer 1982: Berevall Jolaussou 1985: Schombert. Wallin Struck-Mazreell 1990; Mirabel. Lutz Maza 1991: Mirabel. Dottori Lutz 1992: Duc Mirabel 1901. 1998: Duc et al.," This possibility has been recently reinforced by a number of pieces of observational evidence (Schweizer 1982; Bergvall Johansson 1985; Schombert, Wallin Struck-Marcell 1990; Mirabel, Lutz Maza 1991; Mirabel, Dottori Lutz 1992; Duc Mirabel 1994, 1998; Duc et al."443 2000: Yoshida. Taniguchi Muravauia 1991: Braine et al.," 2000; Yoshida, Taniguchi Murayama 1994; Braine et al."444 2000: Weilbacher et al., 2000; Weilbacher et al.445 20003., 2000).446 Also. IIuusbergeer et al. (," Also, Hunsberger et al. ("4471996: 1998) fud an excess of dwarf ealaxies m compact groups of galaxies appareutlv caused x iuteractious amoue eroup moeniboers.,1996; 1998) find an excess of dwarf galaxies in compact groups of galaxies apparently caused by interactions among group members.448 Such formation of idal dwarf galaxies (TDCGs) has also been demonstrated No nunperical simulations of mereieinteracting galaxies (Darues UWernquist 1992: Elincercen. IWaufinan Thomasson 19903).," Such formation of tidal dwarf galaxies (TDGs) has also been demonstrated by numerical simulations of merging/interacting galaxies (Barnes Hernquist 1992; Elmegreen, Kaufman Thomasson 1993)."449 Therefore. tidal formation seems to votentially be an important formation mechanisi for dwarf galaxies (Okazaki Taniguchi 2000 auc references herein).," Therefore, tidal formation seems to potentially be an important formation mechanism for dwarf galaxies (Okazaki Taniguchi 2000 and references therein)."450 One famous tidal debris svstei extends to the northeast of Sevtert’s Sextet (hereafter SS)., One famous tidal debris system extends to the northeast of Seyfert's Sextet (hereafter SS).451 SS is oue of the most zdnous. as well as deusest. compact groups of galaxies (Sevtert 1918a. 1918b: see for a review. Rabacaa 1996).," SS is one of the most famous, as well as densest, compact groups of galaxies (Seyfert 1948a, 1948b; see for a review, Rabaçaa 1996)."452 This group is also a Iflickson compact group (hereafter TCG) of galaxies. WCC 79 (Tickson 1982: 1993).," This group is also a Hickson compact group (hereafter HCG) of galaxies, HCG 79 (Hickson 1982; 1993)."453 Many subsequent studies of SS have iueutioued that the galaxies in SS appear to show morphologically aud dyvuaiicallv peculiar properties (Suleutic Lorre 1983: Rubin. Huuter. Ford 1991: Dettoni Fasano 1993: Mendes de Oliveira Hicksonu 199E: Boufauti et al," Many subsequent studies of SS have mentioned that the galaxies in SS appear to show morphologically and dynamically peculiar properties (Sulentic Lorre 1983; Rubin, Hunter, Ford 1991; Bettoni Fasano 1993; Mendes de Oliveira Hickson 1994; Bonfanti et al."454 1999: Nishiura et al., 1999; Nishiura et al.455 200042)., 2000a).456 IDHickson πιο reearded SS as a galaxw quartet CICG 79a. 79b. 79c. and του).," Hickson himself regarded SS as a galaxy quartet (HCG 79a, 79b, 79c, and 79d)."457 A fifth. component. TCC 79e. was found to be a redshift-discordant galaxy that is believed to have no pliysical relation to SS (Ποσο 1992).," A fifth component, HCG 79e, was found to be a redshift-discordant galaxy that is believed to have no physical relation to SS (Hickson 1992)."458 A sixth object. or more preciselv the north-casteru optical fizz. is now considered likely to be tidal debris associated with the," A sixth object, or more precisely the north-eastern optical fuzz, is now considered likely to be tidal debris associated with the"459ffeatures. since this depends strongly on the brightness. the signal to noise. and the ‘shape’ of the features.,"features, since this depends strongly on the brightness, the signal to noise, and the 'shape' of the features."460 It is hence impossible to derive a general value for the position accuracy of H2--knots., It is hence impossible to derive a general value for the position accuracy of -knots.461 The limiting magnitude 0 detection limit) in the nnarrow band filter ranges mostly between mmag and mmag (bottom panel in refstats))., The limiting magnitude $\sigma$ detection limit) in the narrow band filter ranges mostly between mag and mag (bottom panel in \\ref{stats}) ).462 The limiting magnitude in each image is generally no a result of the sky quality (the vast majority of the data was taken under photometric conditions). but rather is caused by crowdinnr along the Galactic Plane.," The limiting magnitude in each image is generally not a result of the sky quality (the vast majority of the data was taken under photometric conditions), but rather is caused by crowding along the Galactic Plane."463 This is nicely illustrated by the apparen correlation between the limiting magnitude in the images and their position (longitude) along the Galactic Plane (see reflong))., This is nicely illustrated by the apparent correlation between the limiting magnitude in the images and their position (longitude) along the Galactic Plane (see \\ref{long}) ).464 One can see in this figure that for fields more than abou oor ffrom the Galactic Centre. the limiting factor is indeed the integration time. i.e. we reach a point source limit of mmag.," One can see in this figure that for fields more than about or from the Galactic Centre, the limiting factor is indeed the integration time, i.e. we reach a point source limit of mag."465 For fields closer to the Galactic Centre a linear correlation between f and the limiting magnitude is seen. caused by the increased crowding (confusion noise) in the images at these positions.," For fields closer to the Galactic Centre a linear correlation between $l$ and the limiting magnitude is seen, caused by the increased crowding (confusion noise) in the images at these positions."466 For the detection of eemission. which is usually extended. the surface brightness detection limit is of interest.," For the detection of emission, which is usually extended, the surface brightness detection limit is of interest."467 We estimate a typical noise level in our images of 34.110.1° Zaaresec. 7., We estimate a typical noise level in our images of $\cdot$ $^{-19}$ $^{-2}$ $^{-2}$.468 We hence note that a source with a uniform surface brightness of 70 7 > will be detected at a 36 level in our unbinned images ppixel scale}., We hence note that a source with a uniform surface brightness of $^{-18}$ $^{-2}$ $^{-2}$ will be detected at a $\sigma$ level in our unbinned images pixel scale).469 Assuming the ffeatures are extended over several arcseconds. we estimate a 3 detection limit after re-binning our images to e.g. (a typical resolution in GLIMPSE) of 17-110 1? Zaaresee 7.," Assuming the features are extended over several arcseconds, we estimate a $\sigma$ detection limit after re-binning our images to e.g. (a typical resolution in GLIMPSE) of $\cdot$ $^{-19}$ $^{-2}$ $^{-2}$."470 Using the GLIMPSE3¢ point source detection limit spread over the PSF sstr Churchwell et al. (2009)))," Using the GLIMPSE$\sigma$ point source detection limit spread over the PSF $^{-1}$, Churchwell et al. \cite{2009PASP..121..213C}) )"471 and a conversion from fflux at ym into the I-0SSCI) flux for a range of typical conditions (Takami et al. 2010))), and a conversion from flux at $\mu$ m into the S(1) flux for a range of typical conditions (Takami et al. \cite{2010ApJ...720..155T}) )472" one finds that the GLIMPSE 30 detection limit corresponds to 5-40- 1"" = ? in the 1-0SSCI) line of molecular hydrogen.", one finds that the GLIMPSE $\sigma$ detection limit corresponds to $\cdot$ $^{-17}$ $^{-2}$ $^{-2}$ in the S(1) line of molecular hydrogen.473" Our survey is hence a factor of 300 to 2000 better than the corresponding GLIMPSE detections of iin the 4.5,/m filter.", Our survey is hence a factor of 300 to 2000 better than the corresponding GLIMPSE detections of in the $\mu$ m filter.474 Thus. only for regions with a K-band extinction in excess of mmag. might GLIMPSE detect j/im eemission notdetectable in our data.," Thus, only for regions with a K-band extinction in excess of mag, might GLIMPSE detect $\mu$ m emission notdetectable in our data."475 Collectively. jets and outflows can be used as sign-posts of star ormation (e.g. Bally et al. (1995):," Collectively, jets and outflows can be used as sign-posts of star formation (e.g. Bally et al. \cite{1995ApJ...454..345B};"476 Eislófffel 0000: Froebrich Scholz (20033:: Hodapp (20073: Hatehell et al. (2007):, Eislöfffel \cite{2000A&A...354..236E}; Froebrich Scholz \cite{2003A&A...407..207F}; Hodapp \cite{2007AJ....134.2020H}; Hatchell et al. \cite{2007A&A...472..187H};477 Davis et al. (2009):, Davis et al. \cite{2009A&A...496..153D};478 Walawender et al. (2009)))., Walawender et al. \cite{2009AJ....137.3254W}) ).479 An abundance of jets ;»oints to active accretion and a young population: a paucity of flows. in a region where near- and mid-IR photometry still point ο a sizeable population of reddened sources. indicates à more evolved region with a larger population of pre-main-sequence stars CI-Tauri and Herbig Ae/Be stars).," An abundance of jets points to active accretion and a young population; a paucity of flows, in a region where near- and mid-IR photometry still point to a sizeable population of reddened sources, indicates a more evolved region with a larger population of pre-main-sequence stars (T-Tauri and Herbig Ae/Be stars)."480 Indeed. one of the goals of his survey is to establish whether the large number of outflows seen in a massive star forming region like 775 (over 50 independent flows have been detected in a single WFCAM tile Davis et al. (200731). ," Indeed, one of the goals of this survey is to establish whether the large number of outflows seen in a massive star forming region like 75 (over 50 independent flows have been detected in a single WFCAM tile -- Davis et al. \cite{2007MNRAS.374...29D}) ),"481is common in other massive star forming regions., is common in other massive star forming regions.482 Since the dynamical age of a protostellar outflow is 10 to 100-times less than the turbulent lifetime of a GMC. an abundance of outflows in each region would point to ongoing or multiple epochs of star formation. rather than infrequent bursts of star formation in each GMC.," Since the dynamical age of a protostellar outflow is 10 to 100-times less than the turbulent lifetime of a GMC, an abundance of outflows in each region would point to ongoing or multiple epochs of star formation, rather than infrequent bursts of star formation in each GMC."483 Since outflows are also a direct tracer of mass aecretion and ejection they can be used to estimate star formation efficiency. particularly in high mass star forming regions where the efficiency is grossly affected by molecular cores and existing massive voung stars. which influence the environment via their hugely energetic winds and intense UV fluxes.," Since outflows are also a direct tracer of mass accretion and ejection they can be used to estimate star formation efficiency, particularly in high mass star forming regions where the efficiency is grossly affected by molecular cores and existing massive young stars, which influence the environment via their hugely energetic winds and intense UV fluxes."484 An unbiased survey like UWISH2 will map star formation efficiency from region to region. providing a reliable map of the distribution of dynamically active (showing signs of jets/outflows) star forming regions along the Galactic Plane.," An unbiased survey like UWISH2 will map star formation efficiency from region to region, providing a reliable map of the distribution of dynamically active (showing signs of jets/outflows) star forming regions along the Galactic Plane."485 Individually. jets and outflows can be used to pin-point the locations of protostars.," Individually, jets and outflows can be used to pin-point the locations of protostars."486 Flows from 00 sources. for example. are typically |O-times brighter in eemission than their II counterparts (Caratth ο Garatti et al. (2006))).," Flows from 0 sources, for example, are typically 10-times brighter in emission than their I counterparts (Caratti o Garatti et al. \cite{2006A&A...449.1077C}) ),"487 while optically-visible Herbig-Haro jets from III sources are extremely faint in eemission. because of a lack of ambient molecular gas.," while optically-visible Herbig-Haro jets from II sources are extremely faint in emission, because of a lack of ambient molecular gas."488 Outflows from massive young stars are. on the other hand. rarely detected in the optical because of extinction.," Outflows from massive young stars are, on the other hand, rarely detected in the optical because of extinction."489 These flows are usually only seen in eemission (e.g. Varricatt et al. (2010))), These flows are usually only seen in emission (e.g. Varricatt et al. \cite{2010MNRAS.404..661V}) )490 and in millimeter-wave line maps and radio continuum images (Arce et al. (2007)))., and in millimeter-wave line maps and radio continuum images (Arce et al. \cite{2007prpl.conf..245A}) ).491 However. the spatial resolution of the longer wavelength observations is usually poorer.," However, the spatial resolution of the longer wavelength observations is usually poorer."492 With outflow sources identified. the mass of the individual sources. derived from supporting multi-wavelength photometry (WFCAM-GPS. Spitzer. JCMT-GPS. VLA. ete.)," With outflow sources identified, the mass of the individual sources, derived from supporting multi-wavelength photometry (WFCAM-GPS, Spitzer, JCMT-GPS, VLA, etc.)"493 and thus mass/Iuminosity estimates. can be measured.," and thus mass/luminosity estimates, can be measured."494 This sort of analvsis will help to address the question of whether the most massive stars. which may not form through disk accretion. can generate collimated outflows.," This sort of analysis will help to address the question of whether the most massive stars, which may not form through disk accretion, can generate collimated outflows."495 Because fflows are driven by the youngest sources. oobservations can also break the protostar/T-Tauri star ambiguity in Spitzer-IRAC colour-colour analysis. where inclination effects grossly affect mid-IR colours and thereby hinder individual source classifications OO/T or ITE: e.g. Allen et al. 200430).," Because flows are driven by the youngest sources, observations can also break the protostar/T-Tauri star ambiguity in Spitzer-IRAC colour-colour analysis, where inclination effects grossly affect mid-IR colours and thereby hinder individual source classifications 0/I or II; e.g. Allen et al. \cite{2004ApJS..154..363A}) )."496 Jets are also powerful tracers of infall history. since Jet parameters correlate closely with mass infall rates and accretion luminosities (Beck (2007)... Antoniucci et al. (2008))).," Jets are also powerful tracers of infall history, since jet parameters correlate closely with mass infall rates and accretion luminosities (Beck \cite{2007AJ....133.1673B}, Antoniucci et al. \cite{2008A&A...479..503A}) )."497 Tight clustering. interactions between protostellar neighbours. and particularly photo-evaporation and ablation of protostellar disks. can inhibit accretion and thereby switch off the engine that drives an outflow (they can also trigger accretion and cause FU-Ori type outbursts).," Tight clustering, interactions between protostellar neighbours, and particularly photo-evaporation and ablation of protostellar disks, can inhibit accretion and thereby switch off the engine that drives an outflow (they can also trigger accretion and cause FU-Ori type outbursts)."498 This will be particularly important in massive star forming regions. where young stars form in clusters and where massive stars influence their lower-mass neighbours through gravitational. radiative and mechanical (outflow) interactions.," This will be particularly important in massive star forming regions, where young stars form in clusters and where massive stars influence their lower-mass neighbours through gravitational, radiative and mechanical (outflow) interactions."499 Estimates of the frequency of jet activity in clustered environments. provided by UWISH?2. will lead to an assessment of the degree to which interaction inhibits accretion.," Estimates of the frequency of jet activity in clustered environments, provided by UWISH2, will lead to an assessment of the degree to which interaction inhibits accretion."500 Statistical studies of jets can also shed light on the dynamics of cloud collapse and star formation in GMCs., Statistical studies of jets can also shed light on the dynamics of cloud collapse and star formation in GMCs.501 Are outflows randomly orientated. or are they aligned perpendicular to cloud filaments though parallel to magnetic field lines (Eislófffel et al. (1994).. ," Are outflows randomly orientated, or are they aligned perpendicular to cloud filaments though parallel to magnetic field lines (Eislöfffel et al. \cite{1994AJ....108.1042E}, ,"502Banerjee Pudritz (2006)))?, Banerjee Pudritz \cite{2006ApJ...641..949B}) )?503 Existing observations yield contrasting results (e.g. Anathpindika Whitworth (2008 }:: Davis et al. (2009))):, Existing observations yield contrasting results (e.g. Anathpindika Whitworth \cite{2008A&A...487..605A}; ; Davis et al. \cite{2009A&A...496..153D}) );504 clearly. a large. statistically-significant sample of flows. with complementary observations of cloud and magnetic," clearly, a large, statistically-significant sample of flows, with complementary observations of cloud and magnetic"505Each low-degree p-mode of solar-like oscillations carries unique information about stellar interior.,Each low-degree p-mode of solar-like oscillations carries unique information about stellar interior.506 Thus asteroseismology has the capability to probe the interior of stars and to determine the fundamental parameters of individual stars (Ulrich.1986:GoughYang&Meng 20100).," Thus asteroseismology has the capability to probe the interior of stars and to determine the fundamental parameters of individual stars \citep{ulri86, goug87, chri02, egge06, yang10a}."507. However. it is more difficult to extract estimates of individual frequencies than to extract the mean large frequency separation GN) and the frequency of maximum seismic amplitude 9.," However, it is more difficult to extract estimates of individual frequencies than to extract the mean large frequency separation ) and the frequency of maximum seismic amplitude )."508 Fortunately. the and also allow us to determine the fundamental parameters of stars (mass and radius) to within a few percent (Kjeldsen&Bedding1995:Stelloetal. 2009a.b).," Fortunately, the and also allow us to determine the fundamental parameters of stars (mass and radius) to within a few percent \citep{kjel95, stel09a, stel09b}."509. Moreover. the property of frequency separations has been studied by many investigators (Ulrich1986:GoughYang&Meng 2009).," Moreover, the property of frequency separations has been studied by many investigators \citep{ulri86, goug87, gabr89, roxb00, yang07, yang09}."510. Asteroseismic parameters and will be the primary parameters for asteroseismology., Asteroseismic parameters and will be the primary parameters for asteroseismology.511 Using the data observed by the (Baglinetal.2006).. Hekkeretal.(2009) and Mosseretal.(2010) extracted the and of more than 900 red giant stars.," Using the data observed by the \citep{bagl06}, \citet{hekk09} and \citet{moss10} extracted the and of more than 900 red giant stars."512 By making use of the andAv.. Kallingeretal.(2010). and Mosseretal.(2010) derived information on the mass and radius of the oscillating stars.a," By making use of the and, \citet{kall10} and \citet{moss10} derived information on the mass and radius of the oscillating stars.,"513ddition.. the distributions of the and applied to test stellar population synthesis models Tiglioetal.2009:Yang.Meng&Li2010) and physical processes such as mass loss and binary interactions (Yang.Meng&Li2010)..," the distributions of the and applied to test stellar population synthesis models \citep{migl09, yang10b} and physical processes such as mass loss and binary interactions \citep{yang10b}."514 Asteroseismology has significantly advanced the theory of stellar Palructure and evolution., Asteroseismology has significantly advanced the theory of stellar structure and evolution.515 Since stars in a cluster are believed to have uniform age and οjemical compositions. research on solar-like oscillations in cluster oZars as a uniform ensemble will strengthen our ability to Aellar models and improve our understanding of stellar evolution and interior physical processes," Since stars in a cluster are believed to have uniform age and chemical compositions, research on solar-like oscillations in cluster stars as a uniform ensemble will strengthen our ability to stellar models and improve our understanding of stellar evolution and interior physical processes"516for the PDMF assumes a constant star formation rate in the Galactic center.,for the PDMF assumes a constant star formation rate in the Galactic center.517 Our calculations involve the computation of multidimensional integrals over a two dimensional parameter space (see §4))., Our calculations involve the computation of multidimensional integrals over a two dimensional parameter space (see \ref{sec:mass_loss_rates}) ).518" Therefore, for the ease of calculation, we ignore the enhancement of the collision rate due to the effects of gravitational focusing."," Therefore, for the ease of calculation, we ignore the enhancement of the collision rate due to the effects of gravitational focusing."519 This results in a conservative estimate of the collision rate., This results in a conservative estimate of the collision rate.520" As two projectiles collide with each other, their mutual gravitational⋅⋅ attraction⋅ pulls them together, resulting ⋅⋅in an"," As two projectiles collide with each other, their mutual gravitational attraction pulls them together, resulting in an"521bright “spot” poiuted out in (Oppenheimeretal.2008).,bright “spot” pointed out in \citep{Oppe08}.522". The distauces of the major and minor axes of the elliptical fit to the isophotes are indicated by down aud up arrows, respectively."," The distances of the major and minor axes of the elliptical fit to the isophotes are indicated by down and up arrows, respectively."523 We calculate the error bars according to the following procedure., We calculate the error bars according to the following procedure.524" At a given radial distance from the cllipse center, we sample the brightucss aloug the circle of the even radius."," At a given radial distance from the ellipse center, we sample the brightness along the circle of the given radius."525 The inclination of the disk iutroduces au intrinsic variation in brightuess that is svuuetric about only on axis. so we find the best-fitting function of the orm e|bsinte)dsin(20ο) where 0 is the position anele aud calculate the rums deviation from this best-fitting function to be the error.," The inclination of the disk introduces an intrinsic variation in brightness that is symmetric about only on axis, so we find the best-fitting function of the form $a + b\sin(\theta+c) + d\sin(2\theta+e)$ where $\theta$ is the position angle and calculate the rms deviation from this best-fitting function to be the error."526" The errors are simular In magnitude to those caleulated with a hes-fittine ""nuctiou of «|bsin(0e) so we are not over-fitting the data.", The errors are similar in magnitude to those calculated with a best-fitting function of $a + b\sin(\theta+c)$ so we are not over-fitting the data.527 The surface brightuess profiles of the data (Fig., The surface brightness profiles of the data (Fig.528 2 lower oft) are qualitatively simular to those of the plauct-less uodel (Fig., \ref{allprofiles} lower left) are qualitatively simular to those of the planet-less model (Fig.529 2. upper middle). with the steepest profile along the southern minor axis (dotted line}. and the jortheru ninuor axis (solid line) consistently dimuucr than he major axes (dashed and dot-dashed).," \ref{allprofiles} upper middle), with the steepest profile along the southern minor axis (dotted line), and the northern minor axis (solid line) consistently dimmer than the major axes (dashed and dot-dashed)."530 This validates our asstuuption that the nortlowestern edge of the disk is ited away from the observer., This validates our assumption that the northwestern edge of the disk is tilted away from the observer.531 When a plauct is added to the image. it creates a dimming in the surface brightucss profile slightly inward of the leneth of the minor axis aud ai brightening just outward of it.," When a planet is added to the image, it creates a dimming in the surface brightness profile slightly inward of the length of the minor axis and a brightening just outward of it."532 The imaguitude of this S-shaped perturbation erows with ducreasine planet mass., The magnitude of this S-shaped perturbation grows with increasing planet mass.533 The shifts in the profiles along the other axes cau be explained by the shift of the ellipse. center as the planet mass lnereases., The shifts in the profiles along the other axes can be explained by the shift of the ellipse center as the planet mass increases.534 Tn order to establish how well the mocels match the data. we need to quautify the goodness of the fits.," In order to establish how well the models match the data, we need to quantify the goodness of the fits."535 Since we are interested iu the perturbations caused by the xeseuce of a planet and ouly qualitatively interest he rest of the disk structure. we re-scale the brighticccX* xofiles before directly comparing the data to the mockρα," Since we are interested in the perturbations caused by the presence of a planet and only qualitatively interested in the rest of the disk structure, we re-scale the brightness profiles before directly comparing the data to the models."536 For a eiven disk iauage. we first scale the radius to the eusth of the minor or major axis. as appropriate. to derive a normalized distance.," For a given disk image, we first scale the radius to the length of the minor or major axis, as appropriate, to derive a normalized distance."537 In the data nuage. the elliptical radius along the plauct’s position differs frou he nüuor axis bv 14. so it makes little difference to use the nünor axis to scale this profile.," In the data image, the elliptical radius along the planet's position differs from the minor axis by $1\%$, so it makes little difference to use the minor axis to scale this profile."538 We then divide he brightness profile along the planet axis bv that of he major axis., We then divide the brightness profile along the planet axis by that of the major axis.539 This cancels out any small variatious hat nuelt result from uncertainties iu inclinatiou angle., This cancels out any small variations that might result from uncertainties in inclination angle.540 Iu the case of the data image. we average between the xofiles of both major axes before scaling.," In the case of the data image, we average between the profiles of both major axes before scaling."541 The resulting normalized brightness profiles for the data and models are plotted in Fig. 3., The resulting normalized brightness profiles for the data and models are plotted in Fig. \ref{plotratios}.542 Using these ποβ]ος and error bars measured from the data. we calculate the reduced 4? over the normalized radius range Yon 0.6 to 1.3.," Using these profiles and error bars measured from the data, we calculate the reduced $\chi^2$ over the normalized radius range from $0.6$ to $1.3$."543 The resulting values of reduced axe abulated iu Table 1.., The resulting values of reduced $\chi^2_{\nu}$ are tabulated in Table \ref{ellipsetable}.544 We find that both a Saturn 4Zmass Xauet aud no planet at all are eood fits to the data. while a Jupiter mass or larger is excluded.," We find that both a Saturn mass planet and no planet at all are good fits to the data, while a Jupiter mass or larger is excluded."545"wavelengths of Lya andrv, respectively (Fig. 9)),","wavelengths of $\alpha$ and, respectively (Fig. \ref{fig:sepac}) ),"546" M(C,AB) might be slightly wavelength dependent, providing marginal evidence that extinction is lower for C than for AB."," $M$ (C,AB) might be slightly wavelength dependent, providing marginal evidence that extinction is lower for C than for AB."547" Since the differential extinction remains low, we also assume that it does not affect the macro-amplification factor determined at the wavelength of Ha by more than2%."," Since the differential extinction remains low, we also assume that it does not affect the macro-amplification factor determined at the wavelength of $\alpha$ by more than."548". The flux ratios with respect to component A are then F0/Fao = 0.95+0.02 (Fig. 4)),"," The flux ratios with respect to component A are then $F_{\rm549B0}$ $F_{\rm A0}$ =$\pm$ 0.02 (Fig. \ref{fig:ratio}) ),"550 Fco/Fao = 0.93+0.10 and Fpo/Fao = 0.39+0.04.," $F_{\rm551C0}$ $F_{\rm A0}$ = $\pm$ 0.10 and $F_{\rm D0}$ $F_{\rm A0}$ = $\pm$ 0.04."552" The fact that Fco/FA4o and Fpo/FAo are different from the values estimated from photometry is due to a significant de-amplification of the C continuum and to a significant amplification of the D continuum, as derived from the analysis of the spectral lines."," The fact that $F_{\rm C0}$ $F_{\rm A0}$ and $F_{\rm553D0}$ $F_{\rm A0}$ are different from the values estimated from photometry is due to a significant de-amplification of the C continuum and to a significant amplification of the D continuum, as derived from the analysis of the spectral lines."554 This emphasizes the need to properly correct for microlensing before interpreting the flux ratios., This emphasizes the need to properly correct for microlensing before interpreting the flux ratios.555 MacLeod et al. (2009)), MacLeod et al. \cite{mac09}) )556" determined Fpgo/FAo = 0.84+0.07, Foo/Fao = 0.72+0.07 and Fpo/Fao = 0.40+0.06 at 11 μα in the mid-infrared, i.e., where microlensing and extinction are thought to be negligible."," determined $F_{\rm B0} / F_{\rm A0}$ = $\pm$ 0.07, $F_{\rm C0} / F_{\rm A0}$ = $\pm$ 0.07 and $F_{\rm557D0} / F_{\rm A0}$ = $\pm$ 0.06 at 11 $\mu$ m in the mid-infrared, i.e., where microlensing and extinction are thought to be negligible."558" Although only marginally different, the flux ratios of the B and C components relative to A seem slightly smaller than ours."," Although only marginally different, the flux ratios of the B and C components relative to A seem slightly smaller than ours."559" If real, the origin of such a discrepancy is not clear but could be related to the intense starburst activity detected in the host galaxy of ((Lutz et al. 2007,,"," If real, the origin of such a discrepancy is not clear but could be related to the intense starburst activity detected in the host galaxy of (Lutz et al. \cite{lut07}, ,"560" Bradford et al. 2009)),"," Bradford et al. \cite{bra09}) ),"561 which possibly contaminates with PAH emission the 11.2 jum (3.2 um flux measurements., which possibly contaminates with PAH emission the 11.2 $\mu$ m (3.2 $\mu$ m rest-frame) flux measurements.562The lack of multi-epoch observations covering the whole cm range does not allow for conclusiveness.,The lack of multi-epoch observations covering the whole cm range does not allow for conclusiveness.563" ""his source was first studied by VanderVeenctal(1989).. who had selected it as a post-AGB candidate because of its IRAS colours."," This source was first studied by \citet{vanderveen}, who had selected it as a post-AGB candidate because of its IRAS colours."564 Hycdroxil maser lines with Ae~40 km s were detected with the Parkes and Nancaay raciotelescopes bv teLintelHekkertctal(1991). and Szvmezak&CGorard(2004) respectively: the latter detected both 1612 and 1667 Alllz lines., Hydroxil maser lines with $\Delta v\sim40$ km $^{-1}$ were detected with the Parkes and Nançaay radiotelescopes by \citet{telintel} and \citet{szymczak} respectively; the latter detected both 1612 and 1667 MHz lines.565 Nevertheless. VanderVeenetal(1989). did not detect the OL 1612 MllIZ line with the VLAX in D array at the coordinates of the source. but —7 north of it.," Nevertheless, \citet{vanderveen} did not detect the OH 1612 MHz line with the VLA in B array at the coordinates of the source, but $\sim$ $'$ north of it."566 Et is likely then that the detections performed with the single-cish telescopes are to be associated to this oll-set maser. although further interferometric observations would. be desirable to rule out an intrinsic OLL variability.," It is likely then that the detections performed with the single-dish telescopes are to be associated to this off-set maser, although further interferometric observations would be desirable to rule out an intrinsic OH variability."567 Sanchez-Contrerasctal(2004) performed. optical spectroscopy of LRAS 17516-2522525 and found that its spectrum. is dominated by nebular emission. lines., \citet{sanchez08} performed optical spectroscopy of IRAS 17516-2525 and found that its spectrum is dominated by nebular emission lines.568 They also sugeest the presence. of a hot central star. based," They also suggest the presence of a hot central star, based"569realizatioon of the International Celestial Reference Frame (ICRF2).,on of the International Celestial Reference Frame (ICRF2).570The detection by Tytler οἱ al. (,The detection by Tytler et al. (5711995) and by Cowie et al. (,1995) and by Cowie et al. (5721995) of CIV in the Lyman-a forest clouds has provided the first evidence of widespread chemical enrichinent in the IGM al 2~3.,1995) of CIV in the $\alpha$ forest clouds has provided the first evidence of widespread chemical enrichment in the IGM at $z\sim 3$.573 Madau Shull (1996) have computed the ionizing stellar radiation flux which accompanies (he production of metals at high-z. thus realizing that (his mary be significant. comparable to the QSO contribution if a fraction Z25% of the UV radiation emitted from stars can escape into the intergalactic space.," Madau Shull (1996) have computed the ionizing stellar radiation flux which accompanies the production of metals at $z$, thus realizing that this may be significant, comparable to the QSO contribution if a fraction $\gtrsim 25$ of the UV radiation emitted from stars can escape into the intergalactic space."574 As α consecquence. (he importance of the contribution to the UVB by the galaxy population depends on (he ionizing escape Traction of the UV photons which is a poorly known parameter.," As a consequence, the importance of the contribution to the UVB by the galaxy population depends on the ionizing escape fraction of the UV photons which is a poorly known parameter."575 At low redshilts. Giallongo. Fontana. Madan 1997 ancl Shull οἱ al. (," At low redshifts, Giallongo, Fontana, Madau 1997 and Shull et al. ("5761999) have emphasized that the galaxy. contribution is already. comparable to the QSO contribution for an escape fraction of55.,1999) have emphasized that the galaxy contribution is already comparable to the QSO contribution for an escape fraction of.577.. This value is consistent with the upper limits derived [rom the Hopkins Ultraviolet Telescope (IIUT) by Leitherer οἱ al (1995) ancl Iburwitz et al. (, This value is consistent with the upper limits derived from the Hopkins Ultraviolet Telescope (HUT) by Leitherer et al (1995) and Hurwitz et al. (5781997) and more recently from the Far Ultraviolet Spectroscopic Explorer (FUSE) by Deharveng et al. (,1997) and more recently from the Far Ultraviolet Spectroscopic Explorer (FUSE) by Deharveng et al. (5792001).,2001).580 At the intermediate redshifts 2~1 recent preliminary results by Ferguson. (2001) obtained by means of deep HIST. imagine with the Space Telescope Imaging Spectrograph STIS seem to indicate a slightly higher average upper limit for the UV ratio of the order of feSBW., At the intermediate redshifts $z\sim 1$ recent preliminary results by Ferguson (2001) obtained by means of deep HST imaging with the Space Telescope Imaging Spectrograph STIS seem to indicate a slightly higher average upper limit for the UV ratio of the order of $f_{esc} \lesssim 18$.581 All these attempts made al low-intermediate redshilis have failed so far to derive a statistically sienificant detection of the UV escape fraction., All these attempts made at low-intermediate redshifts have failed so far to derive a statistically significant detection of the UV escape fraction.582 Steidel. Pettini Adelberger (SPA) (2001) on the contrary have recently obtained the first significant detection of Lyman limit [lux in a sample of 29 z~3.40--0.09 Lyman break galaxies (LDGs) using spectra obtained in a multi object spectroscopy configuration devoted to the redshift identification of the high z galaxy sample.," Steidel, Pettini Adelberger (SPA) (2001) on the contrary have recently obtained the first significant detection of Lyman limit flux in a sample of 29 $z\simeq 3.40\pm 0.09$ Lyman break galaxies (LBGs) using spectra obtained in a multi object spectroscopy configuration devoted to the redshift identification of the high $z$ galaxy sample."583 The detected Lyman limit flux corresponds to an escape Iraction Of fice=6506..., The detected Lyman limit flux corresponds to an escape fraction of $f_{esc} = 65$.584 A value that is 3.6—13 times higher than previous upper limits., A value that is $3.6 - 13$ times higher than previous upper limits.585 In the same period we have also started a program to measure the UV escape fraction by means ol deep long slit spectra of individual galaxies., In the same period we have also started a program to measure the UV escape fraction by means of deep long slit spectra of individual galaxies.586 This approach avoids possible biases inherent to the procedure of combining spectra of several different galaxies., This approach avoids possible biases inherent to the procedure of combining spectra of several different galaxies.587 We report here the results from two galaxies observed at slightly lower redshift (2~ 3—3.2) with the FORS2 low resolution spectrograph at the ESO Verv Laree Telescope (VLT) telescope in the spectral range 3500—6500.A.," We report here the results from two galaxies observed at slightly lower redshift $z\simeq 3588-3.3$ ) with the FORS2 low resolution spectrograph at the ESO Very Large Telescope (VLT) telescope in the spectral range $3500-6500$."589 Our galaxies are (vo among the bright end of the Steidel οἱ al. (, Our galaxies are two among the bright end of the Steidel et al. (5901996) sample of 2~3 star forming galaxies observable from Paranal.,1996) sample of $z\sim 3$ star forming galaxies observable from Paranal.591 Thev are identified as C2 and D6 and have been, They are identified as C2 and D6 and have been592AMic-intrared fine structure lines are powerful probes of dusty and obscured galactic nuclei: beige able to penetrate extinctions up to the equivalent of Ay~50.,"Mid-infrared fine structure lines are powerful probes of dusty and obscured galactic nuclei, being able to penetrate extinctions up to the equivalent of $_V\sim 50$."593 Using the Short Wavelength Spectrometer (SWS) ou board the Infrared Space Observatory (ISO). it. is )ossible to detect faint Lucs aud sources.," Using the Short Wavelength Spectrometer (SWS) on board the Infrared Space Observatory (ISO), it is possible to detect faint lines and sources."594 The rich observed spectra can be used for a detailed modelling of the ionizins spectra of starbursts (0.9. Rieopoulou et al. 1996..," The rich observed spectra can be used for a detailed modelling of the ionizing spectra of starbursts (e.g. Rigopoulou et al. \cite{rigo96},"595 Kunze et al 1996)), Kunze et al. \cite{kunze96}) )596 aud Ανν (Moorwood ot al. 1996))., and AGNs (Moorwood et al. \cite{moor96}) ).597 Clear differences between their spectra make these lines a valuable new tool for discriminating between ACN and starburst activity in visually obscured ealaxies., Clear differences between their spectra make these lines a valuable new tool for discriminating between AGN and starburst activity in visually obscured galaxies.598 ACN spectra iuclude cussion from highly ionized species aud the so-called coronal lines. requiring plotous up to e»300 00V. for their creation.," AGN spectra include emission from highly ionized species and the so-called coronal lines, requiring photons up to $\sim$ eV for their creation."599 In contrast. starburst spectra are dominated by lines of low excitation species. because even hot. massive stars enit few ionizing xiotous bevoud the edge at 5100V. Line ratios like 225.9;an / 112.spau and L3yan / 12.872. have been used by Lutz al. (1996a))," In contrast, starburst spectra are dominated by lines of low excitation species, because even hot, massive stars emit few ionizing photons beyond the edge at eV. Line ratios like $\mu$ m / $\mu$ m and $\mu$ m / $\mu$ m have been used by Lutz et al. \cite{lutz96a}) )"600 and Ceuzel et al. (1998)), and Genzel et al. \cite{genzel98}) )601 to establish the donnant source of buuinositv iu ultraluniunonus mfrared ealaxies (ULIRGs}., to establish the dominant source of luminosity in ultraluminous infrared galaxies (ULIRGs).602 In some of the starburst templates studied. very faint [OIV| ciission was found. about two orders of magnitude weaker than in typical ACNs.," In some of the starburst templates studied, very faint ] emission was found, about two orders of magnitude weaker than in typical AGNs."603 Faint 1v] cinission im starbursts is not relevant for establishing +1 power source of ULIRGs. but its origi poses au oeteresting problem because its creation ionization energv is shehltly above the edge.," Faint ] emission in starbursts is not relevant for establishing the power source of ULIRGs, but its origin poses an interesting problem because its creation ionization energy is slightly above the edge."604 Iu this letter we examine possible mechauisius for its production., In this letter we examine possible mechanisms for its production.605 Observations ofa variety of starburst galaxies in the 1]. [Nel and [Nerii] lines have been obtained with the ISO- in 1996 aud 1997. as part of a more comprehensive enaranteced time program," Observations of a variety of starburst galaxies in the ], ] and ] lines have been obtained with the ISO-SWS in 1996 and 1997, as part of a more comprehensive guaranteed time program."606 Iu addition. we use data fro a raster of SWS observations along the major axis of obtained on March 16. 1996 in an open time program.," In addition, we use data from a raster of SWS observations along the major axis of obtained on March 16, 1996 in an open time program."607 One of the target lines of this program was [Feu] 25.988yan. which is close enough in wavelength to extract the 1v] ine frou the same scans.," One of the target lines of this program was ] $\mu$ m, which is close enough in wavelength to extract the ] line from the same scans."608 Tuteeration times of our observations in SWSO2 iode were typically 2 seconds per step. ie. 200 seconds or the complete up-down scan coverimg the line.," Integration times of our observations in SWS02 mode were typically 2 seconds per step, i.e. 200 seconds for the complete up-down scan covering the line."609 We used standard procedures from the SWS Iuteractive Analvsis system for data reductiou., We used standard procedures from the SWS Interactive Analysis system for data reduction.610 For strong sources ike bricht starbursts. residual friuging is often obvious in the processed data.," For strong sources like bright starbursts, residual fringing is often obvious in the processed data."611 This was corrected for by fitting sine functions o line-free regious of he spectra., This was corrected for by fitting sine functions to line-free regions of the spectra.612 This xocedure is sufficiently reliable for the brightest sources. since the fringes cau be approximated by a single sine Muction over the small observed wavelength range near Orv}. and the lues are narrow.," This procedure is sufficiently reliable for the brightest sources, since the fringes can be approximated by a single sine function over the small observed wavelength range near ], and the lines are narrow."613 Nevertheless. baseline nucertaity due to fringing is often the limiting factor in measmiug the [Orv| line fiux. and allows us to set only upper limits in some sources. 11]," Nevertheless, baseline uncertainty due to fringing is often the limiting factor in measuring the ] line flux, and allows us to set only upper limits in some sources. ]"614 aud 111] were, and ] were615We plot the average bolometric Luminosity of the extended cluster component (from Table 2)) against quasar redshift in Fieure 4..,We plot the average bolometric luminosity of the extended cluster component (from Table \ref{tab:averages}) ) against quasar redshift in Figure \ref{fig:lxz}.616 For comparison we plot the bolometric luminosity of the X-ray source associated with the distant radio galaxies 3€277.2. BCσοι. 3€C324. BCBH56. BCBGS (Lrom Crawford Fabian 1996b) and 1138-262 (Carilli et al 1998). and the clusters surrounding the two nearby. FR HE radio galaxies Cyenus A (Ueno et al 1994) and 3€295 (Henry Lenriksen 1986).," For comparison we plot the bolometric luminosity of the X-ray source associated with the distant radio galaxies 3C277.2, 3C294, 3C324, 3C356, 3C368 (from Crawford Fabian 1996b) and 1138-262 (Carilli et al 1998), and the clusters surrounding the two nearby FR II radio galaxies Cygnus A (Ueno et al 1994) and 3C295 (Henry Henriksen 1986)."617 The observed countrates of the distant radio galaxies have been converted to luminosities assuming the same tthermal bremsstrahlung mocel used to obtain luminosities for the quasar extended. emission., The observed countrates of the distant radio galaxies have been converted to luminosities assuming the same thermal bremsstrahlung model used to obtain luminosities for the quasar extended emission.618 The luminosities we have derived for the environment of our quasars are brighter than the upper limits of 1.6 10 (rest-DLrame keV)). to. any cluster. emission surrounding three racio-loucl quasars in. Hall ct al (1995. 1997).," The luminosities we have derived for the environment of our quasars are brighter than the upper limits of 1.6 – $\times10^{44}$ (rest-frame ) to any cluster emission surrounding three radio-loud quasars in Hall et al (1995, 1997)."619 We note. however. that those upper limits have been obtained from. images corrected for satellite wobble.," We note, however, that those upper limits have been obtained from images corrected for satellite wobble."620 They also assume therefore that the quasar light follows the stancared LRL PSE derived. bv. David. ct al (1995) and accounts for all the light in the innermost bin., They also assume therefore that the quasar light follows the standard HRI PSF derived by David et al (1995) and accounts for all the light in the innermost bin.621" Our quasar host clusters are consistent with the luminosity of 3.7:10 deteeted by Hall et al (1997) Lor the environment. of the radio-intermediate quasar Η1521|643 at a redshift, >=0.297.", Our quasar host clusters are consistent with the luminosity of $\times10^{45}$ detected by Hall et al (1997) for the environment of the radio-intermediate quasar H1821+643 at a redshift $z$ =0.297.622 The inferred. bolometric luminosities of the extended components we have founc here are completely reasonable for moclerately rich clusters of galaxies at low redshift., The inferred bolometric luminosities of the extended components we have found here are completely reasonable for moderately rich clusters of galaxies at low redshift.623 They are comparable to the luminosities of the clusters associated with the powerful radio galaxies Cvenus A and 3€295., They are comparable to the luminosities of the clusters associated with the powerful radio galaxies Cygnus A and 3C295.624 They are however (Fig., They are however (Fig.625 1) more luminous than the extended X-ray emission detected around more distant. 3CTU radio galaxies above redshift one., 1) more luminous than the extended X-ray emission detected around more distant 3CR radio galaxies above redshift one.626 Whether this indicates evolution. à problem for radio galaxy/quasar unification. or is a result of small number statistics must await the compilation of a complete sample. which are study is not.," Whether this indicates evolution, a problem for radio galaxy/quasar unification, or is a result of small number statistics must await the compilation of a complete sample, which are study is not."627 Alb extended: models have à. central. cooling time considerably shorter than a Hubble. time., All extended models have a central cooling time considerably shorter than a Hubble time.628 We therefore explore the properties of the implied cooling Lows occurring around. these quasars by deriving some approximate parameters from the broken power-law fits to the profiles., We therefore explore the properties of the implied cooling flows occurring around these quasars by deriving some approximate parameters from the broken power-law fits to the profiles.629 We attribute all the X-ray luminosity o£ the extended component within radius 2 to thermal bremsstrahlung from gas with electron density ο (where nxr 1) at a temperature of keV., We attribute all the X-ray luminosity of the extended component within radius $R$ to thermal bremsstrahlung from gas with electron density $n$ (where $n\propto r^{-1}$ ) at a temperature of .630. The cooling time of the gas at /? (except in the case of 9035.1) is then between about 13 billion vr., The cooling time of the gas at $R$ (except in the case of 3C275.1) is then between about 1--3 billion yr.631" We then estimate the cooling How radius rcp at which the cooling time is 10Lu vr and obtainH a rough indication of the mass deposition rate within that racius [rom⋅ the ratio ""Mof the mass of⋅ gas within""M rep to 1077, vr.", We then estimate the cooling flow radius $r_{\rm CF}$ at which the cooling time is $10^{10}$ yr and obtain a rough indication of the mass deposition rate within that radius from the ratio of the mass of gas within $r_{\rm CF}$ to $10^{10}$ yr.632 MThe derived values are shown in Table 3.., The derived values are shown in Table \ref{tab:cf}.633 Note they are of course subject to not only the appropriateness of the fixed. slopes chosen for our original broken power-law model. but also to the true gravitational potential of any cluster. ancl the amount of gravitational work done on the cooling gas.," Note they are of course subject to not only the appropriateness of the fixed slopes chosen for our original broken power-law model, but also to the true gravitational potential of any cluster, and the amount of gravitational work done on the cooling gas."634 The values should be regarded as uncertain by a least a factor of 2., The values should be regarded as uncertain by a least a factor of 2.635 Phey may be underestimated by a factor of at least 2 if the eas tempcratures are significantly higher than the 4 keV assumed and there is internal absorption such as is common in low redshift cooling Lows., They may be underestimated by a factor of at least 2 if the gas temperatures are significantly higher than the 4 keV assumed and there is internal absorption such as is common in low redshift cooling flows.636 Note that. the radius of the surface brightness break which we infer is in the range 40.90 κρο and is similar to the break radius in the profile of the cluster around LRAS 09104 (~60 kpe. Crawford Fabian 1995b).," Note that the radius of the surface brightness break which we infer is in the range 40–90 kpc, and is similar to the break radius in the profile of the cluster around IRAS 09104 $\sim60$ kpc, Crawford Fabian 1995b)."637" ""This is likely to be the radius of the core of the gravitational potential of the cluster: the r3 profile then occurs. within there since that gas is cooling at approximately constant pressure.", This is likely to be the radius of the core of the gravitational potential of the cluster; the $r^{-1}$ profile then occurs within there since that gas is cooling at approximately constant pressure.638 Such small eravitational core radii are characteristic of relaxed lensing cluster cores such as are associated with massive cooling lows (Allen 1998)., Such small gravitational core radii are characteristic of relaxed lensing cluster cores such as are associated with massive cooling flows (Allen 1998).639 Phe large break racii found for 3€275.1 and I11821|643 do not agree with this picture ancl require more detailed images., The large break radii found for 3C275.1 and H1821+643 do not agree with this picture and require more detailed images.640 We can also use our cooling Dow. paranieters to derive a gas pressure 2? at a radius of ((see Table 3)) for comparison to the pressures derived from the completely. independent method. using the ionization state of the extended optical emission lines (Crawford Fabian 1989: Bremer et al 1992: Crawford Vanelerreist 1997)., We can also use our cooling flow parameters to derive a gas pressure $P$ at a radius of (see Table \ref{tab:cf}) ) for comparison to the pressures derived from the completely independent method using the ionization state of the extended optical emission lines (Crawford Fabian 1989; Bremer et al 1992; Crawford Vanderreist 1997).641 The pressures derived. from. the optical nebulosities mostly underestimate those derived from the N-ray. profile fits (Figure 5)) by a factor of up to 10., The pressures derived from the optical nebulosities mostly underestimate those derived from the X-ray profile fits (Figure \ref{fig:pvsp}) ) by a factor of up to 10.642 Given that the gas pressures are mostly derived. from the optical nebulosity using a conservative underestimate to the crucial but unknown UV and soft. X-ray band. of the ionizing nuclear μα»ectrum. this cliscrepaney is not surprising.," Given that the gas pressures are mostly derived from the optical nebulosity using a conservative underestimate to the crucial but unknown UV and soft X-ray band of the ionizing nuclear spectrum, this discrepancy is not surprising."643 Support for js interpretation of the disparitv in derived: pressures: is =ound in Crawford et al (1991) where a better knowledge of 1ο lonizing continuum of the nucleus of 3€263 was found to increase the optically-cerivecl pressure by up to an order of magnitude., Support for this interpretation of the disparity in derived pressures is found in Crawford et al (1991) where a better knowledge of the ionizing continuum of the nucleus of 3C263 was found to increase the optically-derived pressure by up to an order of magnitude.644 In addition. we note that 3€254. the quasar with 10 best agreement between the two pressure values is the μην one where UV LIST data has been used to constrain the garape of the ionizing continuum (Crawford Vanelerricst 1997).," In addition, we note that 3C254, the quasar with the best agreement between the two pressure values is the only one where UV HST data has been used to constrain the shape of the ionizing continuum (Crawford Vanderriest 1997)."645 We note that the derived cluster Iuminosity for οςστὸ is high and implies the presence of a rich cluster which is not seen at other wavelengths., We note that the derived cluster luminosity for 3C273 is high and implies the presence of a rich cluster which is not seen at other wavelengths.646" At its relatively low redshift, of z=0.16. such a cluster should be obvious in he optical band."," At its relatively low redshift of $z=0.16$, such a cluster should be obvious in the optical band."647 As mentioned already. it has the profile most susceptible to svstematic errors. and the absence of an optical cluster may argue for the existence of such errors.," As mentioned already, it has the profile most susceptible to systematic errors, and the absence of an optical cluster may argue for the existence of such errors."648 Lf he PSF is then uncertain by a relative amount equal to the observed 36273 profile and our empirical PSE. then it wil not change greatly. our results on the other quasars excep »rhaps for 3C48.," If the PSF is then uncertain by a relative amount equal to the observed 3C273 profile and our empirical PSF, then it will not change greatly our results on the other quasars except perhaps for 3C48."649 ]nverse-Compton scattering of quasar radiation. coul still contribute to. our. extended. component of X-ray emission., Inverse-Compton scattering of quasar radiation could still contribute to our extended component of X-ray emission.650 Ifa significant process. the X-ray source. woule appear asvmametric ancl lop-sided. as (back) scattering by," Ifa significant process, the X-ray source would appear asymmetric and lop-sided, as (back) scattering by"651purely on redshift selection relative to ITE and metal features.,purely on redshift selection relative to HI and metal features.652 There are seven new candidates with detectable deuterimm features at the melt redshift. iu addition to the previously published candidates from Ruecrs aud Hosau (1996ah).," There are seven new candidates with detectable deuterium features at the right redshift, in addition to the previously published candidates from Rugers and Hogan (1996ab)."653 The reason that they have not been published is that they are nof very convincing. or do not vield a precise abundance.," The reason that they have not been published is that they are not very convincing, or do not yield a precise abundance."654 Oue new example is shown iu figure 1., One new example is shown in figure 1.655 The fitted Doppler parameters (figure 2) show that the claimed deuterium detections are at least mostly real. even if individual cases are suspect.," The fitted Doppler parameters (figure 2) show that the claimed deuterium detections are at least mostly real, even if individual cases are suspect."656 This is true both for the new candidates aud for the previously published oues., This is true both for the new candidates and for the previously published ones.657 The D features have statistically narrower profiles than the preteud candidates. consistent with the deuterium ideutification.," The $D$ features have statistically narrower profiles than the pretend candidates, consistent with the deuterium identification."658 Since we believe the deuterium is statistically real. we can derive a statistical abundance.," Since we believe the deuterium is statistically real, we can derive a statistical abundance."659" For the new sample of absorbers. (that is. the uniform sample exeludiug previously published values). we obtain (Ιου)ΠΠ=3.7540,51."," For the new sample of absorbers, (that is, the uniform sample excluding previously published values), we obtain $\langle \log (D/H)\rangle= -3.75\pm 0.51$."660 The: reduced 47 241is 0.70.sp) indicatingnoH: consistency: with: a universal: abuudance.," The reduced $\chi^2$ is 0.70, indicating consistency with a universal abundance."661 Thisqu is certainly not the case for the sample of pretend absorbers: even confuius ourselves to those with high IIT column (so as to be mere directly comparable to the real ones). the abundance of ρουια is dog(P//T);=3.24 2.6. with a reduced \? of 8.7.," This is certainly not the case for the sample of pretend absorbers; even confining ourselves to those with high HI column (so as to be more directly comparable to the real ones), the abundance of “pretendium” is $\langle \log (P/H)\rangle= -3.2\pm 2.6$ , with a reduced $\chi^2$ of 8.7."662 Other statistics also display differences;, Other statistics also display differences.663 While the real deuterim candidates display cousisteucyv with a universal abundance (linear regression coefficient between VCD) aud NCL) of 1.00). the pretend sample is a scatter plot (linear regression cocficient between NCP?) aud NP) of 0.55.)," While the real deuterium candidates display consistency with a universal abundance (linear regression coefficient between $N(D)$ and $N(H)$ of 1.00), the pretend sample is a scatter plot (linear regression coefficient between $N(P)$ and $N(H)$ of $0.55$ .)"664 These statistics are reflected in the scatter plot shown in figure 3., These statistics are reflected in the scatter plot shown in figure 3.665 Aud the cistributions of N(D) aud N(P) also ciffor. in the sense that P is not as cobunon as D — another wav of saving that for the most part. iuterlopers may be there but statistically have a lower column density than the real deuteriuu.," And the distributions of $N(D)$ and $N(P)$ also differ, in the sense that $P$ is not as common as $D$ — another way of saying that for the most part, interlopers may be there but statistically have a lower column density than the real deuterium."666" Although deuterimm is detected. its primordial abuucdanuce is still uncertain by an order of maenitude,"," Although deuterium is detected, its primordial abundance is still uncertain by an order of magnitude."667 At prescut. we have a very firm lower limit D/II> on primordial deuteritm. from Galactic measurements (Linsky ct al.," At present, we have a very firm lower limit $D/H\ge 2\times 10^{-5}$ on primordial deuterium, from Galactic measurements (Linsky et al."668 1993. 1996). ax well as from some quasar absorbers (c.g. Tytler et al.," 1993, 1996), as well as from some quasar absorbers (e.g. Tytler et al."669 1996)., 1996).670 There is some evidence for a somewhat hieher lower lumit (D/IT>lL.10 >) from quasar absorbers (Sougaila et al., There is some evidence for a somewhat higher lower limit $D/H\ge 4\times 10^{-5}$ ) from quasar absorbers (Songaila et al.671 1996)., 1996).672 A higher abundance than this is not clearly required by the data. although there is some statistical evidence for it and there is no very strong evidence against it. since the low abundanuces are still found in just a few cases where deuterimm may have been destroved.," A higher abundance than this is not clearly required by the data, although there is some statistical evidence for it and there is no very strong evidence against it, since the low abundances are still found in just a few cases where deuterium may have been destroyed."673 Better even than a iach larger statistical study would be to fiud a system wherea clear signature of deuterim cau be proven. or where the iuterloper," Better even than a much larger statistical study would be to find a system wherea clear signature of deuterium can be proven, or where the interloper"674the tidal forces acting on the dark matter.,the tidal forces acting on the dark matter.675" We do not attempt to model the stripping of the dark matter, as there are already several published analytic studies which reproduce the dark matter stripping and tidal heating in simulations well (e.g., Taylor Babul 2001; Benson et 22002)."," We do not attempt to model the stripping of the dark matter, as there are already several published analytic studies which reproduce the dark matter stripping and tidal heating in simulations well (e.g., Taylor Babul 2001; Benson et 2002)."676" Instead, the analytic model proposed in §3.1 simply assumes that, within the stripping radius, the properties of the galaxy are unchanged from their initial state."," Instead, the analytic model proposed in 3.1 simply assumes that, within the stripping radius, the properties of the galaxy are unchanged from their initial state."677" Thus, the dark matter halo is assumed to maintain its initial NFW configuration within this radius."," Thus, the dark matter halo is assumed to maintain its initial NFW configuration within this radius."678" In Appendix B, we present a simple analytic argument that validates this assumption for systems where the mass of the galaxy is less than about of the mass of the group."," In Appendix B, we present a simple analytic argument that validates this assumption for systems where the mass of the galaxy is less than about of the mass of the group."679" We have also directly computed the evolution of the tidal radius (r;, defined in Binney Tremaine 1987; see also Appendix B) of the galaxy in the simulations as a function of time."," We have also directly computed the evolution of the tidal radius $r_t$, defined in Binney Tremaine 1987; see also Appendix B) of the galaxy in the simulations as a function of time."680 In 77 we compare the tidal radius with the radial extent of the hot gaseous halo., In 7 we compare the tidal radius with the radial extent of the hot gaseous halo.681 The tidal radius shrinks at pericentre and then expands but at all times is safely larger than the gaseous halo by at least a factor of 2., The tidal radius shrinks at pericentre and then expands but at all times is safely larger than the gaseous halo by at least a factor of 2.682 In 66 we examined the ram pressure stripping of a galaxy on the most common orbit seen in cosmological simulations., In 6 we examined the ram pressure stripping of a galaxy on the most common orbit seen in cosmological simulations.683 We now experiment with varying the initial orbital parameters., We now experiment with varying the initial orbital parameters.684 This will have the effect of changing both the shape and normalisation of P:am(t)., This will have the effect of changing both the shape and normalisation of $P_{\rm ram}(t)$.685 We use 22 of Benson (2005) to select a range of cosmologically likely orbits; the initial velocity of some orbits is dominated by the radial component while others have nearly circular motionsinitially?., We use 2 of Benson (2005) to select a range of cosmologically likely orbits; the initial velocity of some orbits is dominated by the radial component while others have nearly circular motions.686. We plot the mass loss curves for six such orbits in 88., We plot the mass loss curves for six such orbits in 8.687 The mass loss curves in 88 exhibit a variety of behaviours., The mass loss curves in 8 exhibit a variety of behaviours.688 Orbits that initially have a significant tangential component (and have a total velocity of ~ vc) typically undergo only one pericentric passage over the course of 10 Gyr., Orbits that initially have a significant tangential component (and have a total velocity of $\sim v_c$ ) typically undergo only one pericentric passage over the course of 10 Gyr.689" Consequently, their associated M(t) curves tend to"," Consequently, their associated $M(t)$ curves tend to"690generation takes places in two strong shocks at the stream boundaries (clearly. seen in the V:e panel).,generation takes places in two strong shocks at the stream boundaries (clearly seen in the $\bmath \nabla\cdot \bmath v$ panel).691 All of the stream matter has to pass through the intersection of these shock structures. where it essentially attains its final entropy and becomes comparable to the entropy. of the ambient material. at least in the case of the logP/p.," All of the stream matter has to pass through the intersection of these shock structures, where it essentially attains its final entropy and becomes comparable to the entropy of the ambient material, at least in the case of the $\log P/\rho^{\gamma}$."692 For the entropy the dilference between ambient and stream matter is largely caused by the cillerence in molecular weights and its subsequent. change is due to mixing of these two mecdia rather than dissipation., For the entropy the difference between ambient and stream matter is largely caused by the difference in molecular weights and its subsequent change is due to mixing of these two media rather than dissipation.693" At the point of deepest penetration (2102"" em). the stream matter has a temperature. of G«10' IX and a density of ~3gem7. high enough for nuclear burning (see 54)."," At the point of deepest penetration $\sim 2\times69410^{10}\,$ cm), the stream matter has a temperature of $6 \times 10^7\,$ K and a density of $\sim 3 \gcm$, high enough for nuclear burning (see 5.4)."695 The Ix-HE instability affects both the stream shape and the stream width., The K-H instability affects both the stream shape and the stream width.696 For small relative velocities. the stream is “sausage-shaped and becomes increasingly “snake”-shaped as the relative velocity increases (see Fig. 3)).," For small relative velocities, the stream is “sausage”-shaped and becomes increasingly “snake”-shaped as the relative velocity increases (see Fig. \ref{fig:serp}) )."697 It becomes more important as the stream narrows. since perturbations of a given wavelength: are more important for a narrower How.," It becomes more important as the stream narrows, since perturbations of a given wavelength are more important for a narrower flow."698 LH therefore determines the minimum stream width and thus limits the depth of penetration., It therefore determines the minimum stream width and thus limits the depth of penetration.699 In particular. à stream that is very narrow initially will penetrate less deep than a stream with a similar parameters like velocity. density and pressure. but with a larger mass-Iow rate and hence Larger initial stream width.," In particular, a stream that is very narrow initially will penetrate less deep than a stream with a similar parameters like velocity, density and pressure, but with a larger mass-flow rate and hence larger initial stream width."700 In some cases. in particular near the ene of the mereing process. the envelope surrounding the spiraling-in components nav no longer remain in co-rotation with the binary (either because the spiralin time-scale becomes too short or the region surrounding the primary core expands rapidlv).," In some cases, in particular near the end of the merging process, the envelope surrounding the spiraling-in components may no longer remain in co-rotation with the binary (either because the spiral-in time-scale becomes too short or the region surrounding the primary core expands rapidly)."701 Then the stream feels an additional force. due to the moving external medium., Then the stream feels an additional force due to the moving external medium.702 This allects the streams trajectory. in particular the angle of incidence. and thereby 10 penetration depth.," This affects the stream's trajectory, in particular the angle of incidence, and thereby the penetration depth."703 Vhis is illustrated in Figure 4. which galows wo stream calculations with the same parameters as in Figure 2. but where the medium is assumed to rotate with an angular velocity Quan=10tracks relative to the frame of the binary (rotating with angular velocity ©) either in the forward or in the backward direction.," This is illustrated in Figure \ref{fig:rot} which shows two stream calculations with the same parameters as in Figure \ref{fig:pr_disc}704 but where the medium is assumed to rotate with an angular velocity $\Omega_{\rm shift} = 10^{-4}\rads$ relative to the frame of the binary (rotating with angular velocity $\Omega$ ) either in the forward or in the backward direction."705" The frame of the medium (in which the calculation is performed) then rotates with a velocity Quang,=2+Ou."," The frame of the medium (in which the calculation is performed) then rotates with a velocity $\Omega_{\rm706amb} = \Omega \pm \Omega_{\rm shift}$."707 In this frame. he stream itself rotates either anti-clockwise (top panel) or clockwise (bottom panel).," In this frame, the stream itself rotates either anti-clockwise (top panel) or clockwise (bottom panel)."708 As the stream is being pushed xilewards by ambient matter moving against its direction of rotation. Hs trajectory steepens ancl consequently. the penetration depth increases (top panel). while the opposite lappens when the stream is pushed from. behind. (bottom xuwiel).," As the stream is being pushed backwards by ambient matter moving against its direction of rotation, its trajectory steepens and consequently the penetration depth increases (top panel), while the opposite happens when the stream is pushed from behind (bottom panel)."709" As a consequence the stream. penetrates deeper to : ⋅ -Lu ⋜↧↓⋅⋯⊔⊔⊳∖∪⇂∿↓⋅⋀⋅↱≻↓∪≼∼⊔↓↓⊔⊳∖↿∢⊾⋯⇂∪⇂∿−≽ : ⋅−10?"" em in ↓↕∢⋅⊳∖∣⋜⋯∠⇂⋜⊔⋅∠⇂≼∼∥⊳∖∢⊾↿∖↓⊲↝↓⋏∙≟⊳⇉∩⊳∖∖⊽↓↥∐⋖⋅"," As a consequence the stream penetrates deeper to a radius of $\sim7101.75\times 10^{10}\,$ cm instead of $\sim 2\times 10^{10}\,$ cm in the standard case (Fig. \ref{fig:pr_disc}) ),"711⊲↓↿∪⊔↓∙∖⇁↓⋅∢⋅⋯⇍↓↥∢⊾⊳∖⋜⊔⇂⋖⊾↓≻↿↥∪⇂⋅ 5.2.25510 em in. the opposite. case.," while it only reaches a depth of $\sim 2.25\times 10^{10}\,$ cm in the opposite case."712 opThe corresponding. change in the stream temperature at the deepest. point is only a few per cent., The corresponding change in the stream temperature at the deepest point is only a few per cent.713 Note also that for the stream rotating anti-clockwise. the entropy generation at the back side of the stream. where the dynamical pressure is smaller than at the front side (ancl hence has a larger jump in pressure). is largerὃν than for stationary inclined stream.," Note also that for the stream rotating anti-clockwise, the entropy generation at the back side of the stream, where the dynamical pressure is smaller than at the front side (and hence has a larger jump in pressure), is larger than for stationary inclined stream."714 As the stream interacts with the core. it injects not only matter but also angular momentum into the core. spinning," As the stream interacts with the core, it injects not only matter but also angular momentum into the core, spinning"715stars is likely to provide a more accurate description of this galaxy.,stars is likely to provide a more accurate description of this galaxy.716 Past studies of PIX8S2250-41. have suggested. that the prominent emission line arc to the west of the host galaxy originates [rom a direct. collision between the radio source jet and a companion galaxy (Clark ct al 1997. VM99).," Past studies of PKS2250-41 have suggested that the prominent emission line arc to the west of the host galaxy originates from a direct collision between the radio source jet and a companion galaxy (Clark et al 1997, VM99)."717 Figure 11. displavs the infrared. and optical continuum and the ΟΠΗ) emission line imaging data for PI&S2250-41 overlaid with radio contours (all data except. infrared from T05)., Figure \ref{overlays} displays the infrared and optical continuum and the ] emission line imaging data for PKS2250-41 overlaid with radio contours (all data except infrared from T05).718 Continuum emission 1s detected. in both the Avs and E547M filters within the arc. coincident with the secondary radio hotspot.," Continuum emission is detected in both the $K_S$ and F547M filters within the arc, coincident with the secondary radio hotspot."719 We have determined the [ux density of this feature in both filters: the resulting ratio of As emission is 24.3458%., We have determined the flux density of this feature in both filters; the resulting ratio of $K_S$ emission is $24.3 \pm 58\%$.720 This value is consistent with unrecdened stellar. populations aged between either 0.006-0.009Cvr or 0.05-0.1€vr. or alternatively redcdened YSPs of even vounger ages.," This value is consistent with unreddened stellar populations aged between either 0.006-0.009Gyr or 0.05-0.1Gyr, or alternatively reddened YSPs of even younger ages."721 Lor many vears now it has been known that radio source jets have the potential to substantially allect the ambient interstellar/intergalactic medium., For many years now it has been known that radio source jets have the potential to substantially affect the ambient interstellar/intergalactic medium.722" Separating the influence of the jet from other forms of the alignment. elfect (Chambers. Miley van Breugel LOST. AleCarthy et al 1987). Le. scattered AGN emission or line ancl nebular continuum enission associated with ACXN-ionized clouds. is not always straightforward (c.g. Tadhunter et al 2002: Enskip et al 2005) but in many cases there is clear-cut evidence for so-called ""jet.cloud interactions’."," Separating the influence of the jet from other forms of the alignment effect (Chambers, Miley van Breugel 1987, McCarthy et al 1987), i.e. scattered AGN emission or line and nebular continuum emission associated with AGN-ionized clouds, is not always straightforward (e.g. Tadhunter et al 2002; Inskip et al 2005) but in many cases there is clear-cut evidence for so-called `jet–cloud interactions'."723 This may take the form of shocks associated with interactions between the radio jet ancl cool gas clouds in the ISAL/LGAL. or alternatively. jet-inclucecl star formation.," This may take the form of shocks associated with interactions between the radio jet and cool gas clouds in the ISM/IGM, or alternatively, jet-induced star formation."724 Jet-induced star formation was first. proposed in order to explain the observations of such sources as Centaurus A (Blanco ot al 1975: Graham Price 1951: Broclie. WWonniel Bowyer 1983: Sutherland. Bicknell Dopita 1993: Schiminovich et al 1994: Oosterloo Morganti 2005) ancl Minkowski's Object (van Breugel et al 1985: Croft. et al 2006).," Jet-induced star formation was first proposed in order to explain the observations of such sources as Centaurus A (Blanco et al 1975; Graham Price 1981; Brodie, Könnigl Bowyer 1983; Sutherland, Bicknell Dopita 1993; Schiminovich et al 1994; Oosterloo Morganti 2005) and Minkowski's Object (van Breugel et al 1985; Croft et al 2006)."725 Current theoretical mociels (e.g. Fragile ct al 2004: Aellema. Ixurk Rotttecring 2002. and references therein) sugeest that radio source shocks propagating through the clumpy ISMIGAM trigger the collapse and/or fragmentation of overdense regions. which may then subsequently. form stars.," Current theoretical models (e.g. Fragile et al 2004; Mellema, Kurk Rötttgering 2002, and references therein) suggest that radio source shocks propagating through the clumpy ISM/IGM trigger the collapse and/or fragmentation of overdense regions, which may then subsequently form stars."726 Good. evidence for jet-induced. star. formation has also been observed in the case of the radio sources 4€ 41.17 (Dev ct al 1997: Bicknell et al 2000). 3€34. (Best. Longair Robtttecring 1997) and 3€48 (Stockton et al 2007). but on the whole. while evidence for shocks associated with the radio source is relatively common. jetinduced," Good evidence for jet-induced star formation has also been observed in the case of the radio sources 4C 41.17 (Dey et al 1997; Bicknell et al 2000), 3C34 (Best, Longair Rötttgering 1997) and 3C48 (Stockton et al 2007), but on the whole, while evidence for shocks associated with the radio source is relatively common, jet-induced"727are eiven here.,are given here.728 Indeed. the entire range of observed image motion is about 0.2 pixels. and (his is dominated by (he spacecraft guiding precision.," Indeed, the entire range of observed image motion is about 0.2 pixels, and this is dominated by the spacecraft guiding precision."729 Various authors. including Ning(1983).ancl Ixaiseretal.(2000).. have developed theoretical expectations for the astrometric precision of an image.," Various authors, including \citet{kin83} and \citet{kai00}, have developed theoretical expectations for the astrometric precision of an image."730 A simple approximation ls where FWIIM is the image full width at half maximum. and SNR is the photometric signal-to-noise ratio of that star image.," A simple approximation is where FWHM is the image full width at half maximum, and SNR is the photometric signal-to-noise ratio of that star image."731 The differences in the theoretical derivations concern the exact value lor which the approximate value of 2 is used above., The differences in the theoretical derivations concern the exact value for which the approximate value of 2 is used above.732 The observational confirmation of this relationship has vet to be done. but essentially all groumd-. and space-based. astrometric studies have demonstrated (he validity of the scaling of this relationship.," The observational confirmation of this relationship has yet to be done, but essentially all ground- and space-based astrometric studies have demonstrated the validity of the scaling of this relationship."733 Luproved astrometric precision is obtained for smaller image EWILIM. higher ONR. or both assuming that adequate image sampling is available.," Improved astrometric precision is obtained for smaller image FWHM, higher SNR, or both assuming that adequate image sampling is available."734Kepler operates in a heretofore unstudied astrometric domain., operates in a heretofore unstudied astrometric domain.735 The pixels are very large. 3.98 arcseconds. as compared to other ground- and space-basecl astrometric assets. and (he observed FWIIM is approximately 5 to 6 arcseconds and depends on the location in the field of view. (," The pixels are very large, 3.98 arcseconds, as compared to other ground- and space-based astrometric assets, and the observed FWHM is approximately 5 to 6 arcseconds and depends on the location in the field of view. ("736See Drvsonetal.(2010). for further discussion and examples.),See \citet{bry10} for further discussion and examples.)737 The effects of undersampled image components are not captured by Eq., The effects of undersampled image components are not captured by Eq.738 1., 1.739 However.Kepler was designed lor extremely. high. SNR. observations.," However, was designed for extremely high SNR observations."740 The well capacity of the 27-micron CCD pixels is more than a million electrons. and most of the stars are bright.," The well capacity of the 27-micron CCD pixels is more than a million electrons, and most of the stars are bright."741 As more fully diseussed by Caldwellatal.(2010).. the onset of saturation in the basic 6.02 second integration. cvcle is. near the magnitude. Ap = 11.3..," As more fully discussed by \citet{cal10}, the onset of saturation in the basic 6.02 second integration cycle is near the magnitude = 11.3."742(tinge patter: corresponds to Fizeau-like [riuges. the parasitic interference. induces two further coimnponelnts τς» the sigual: Fizeau-like mirror fringes aud Youug-like [riuges.,"fringe pattern corresponds to Fizeau-like fringes, the parasitic interference induces two further components to the signal: Fizeau-like mirror fringes and Young-like fringes."743 The latter coutributiou is incdepeuclen of the object position in the sky., The latter contribution is independent of the object position in the sky.744 We will develop this description in the This interference corrupts the value of interferometric observables: modulus aud phase of the complex. visibilitv., We will develop this description in the This interference corrupts the value of interferometric observables: modulus and phase of the complex visibility.745" The uuderstaucing of this phenomenon is of general interest. for optical long baseline iutererometers which combine multiple beams aud often require some ""compactuess"" iu the opto-mechanical In Section 2. a simple formalism is developed. explicitly showing the different. contributions ‘the parasitic iterfereuce."," The understanding of this phenomenon is of general interest for optical long baseline interferometers which combine multiple beams and often require some ""compactness"" in the opto-mechanical In Section 2, a simple formalism is developed, explicitly showing the different contributions of the parasitic interference."746 In Section 3. the resulting interferometric observables are written for le general case of an exteuded source.," In Section 3, the resulting interferometric observables are written for the general case of an extended source."747 Then the theoretical cases of an uuresolved source and . à stellar systeu with a hot Jupiter are highlighted., Then the theoretical cases of an unresolved source and of a stellar system with a hot Jupiter are highlighted.748 Sections [| aud 5 quantify the impact. of le parasitic inte‘ference ou the theoretical phase signal of these types of sources., Sections 4 and 5 quantify the impact of the parasitic interference on the theoretical phase signal of these types of sources.749 The parameters volved iu this €uantitative study are the factor of parasitic flux between beams. the piston. aud ilje photometric inibalance between interferometric arius.," The parameters involved in this quantitative study are the factor of parasitic flux between beams, the piston, and the photometric imbalance between interferometric arms."750 The impact of the parasitic interference [9]1 hot Jupiter detection is evaluated by using dilfereut syuthetie spectra of such extrasolar plauets., The impact of the parasitic interference on hot Jupiter detection is evaluated by using different synthetic spectra of such extrasolar planets.751 lequiremeuts on the straylieht level limits are also given., Requirements on the straylight level limits are also given.752 Iu general. long-baseliue stellar interferonietry. cousists of sampling au iucident wave packet N —jeans of telescopes at clifferent locatious.," In general, long-baseline stellar interferometry consists of sampling an incident wave packet by means of telescopes at different locations."753 These coherent waves are combiued aud tle resultine[n] interlereuce pattern is extracted in order to measure the complex degree of coherence of the radiatiou jeld., These coherent waves are combined and the resulting interference pattern is extracted in order to measure the complex degree of coherence of the radiation field.754 Then the brightuess distributiou of the source cau be recovered thauks to the Vau Cittert aud Zeruike theorem., Then the brightness distribution of the source can be recovered thanks to the Van Cittert and Zernike theorem.755 This complex degree of coliereuce corresponds to the covariauce of the electric ields collected bv the telescopes., This complex degree of coherence corresponds to the covariance of the electric fields collected by the telescopes.756 If a parasitic interference occurs between the collection aud he recombination steps. the “intrinsic” colerence between beams aud cousequeutly the resulting interferometric ol;ervables will be To our knowledge. the issue of parasitic interference has never been formalizecl except for the digital wave-[rout measuring interferometry technique.," If a parasitic interference occurs between the collection and the recombination steps, the 'intrinsic' coherence between beams and consequently the resulting interferometric observables will be To our knowledge, the issue of parasitic interference has never been formalized except for the digital wave-front measuring interferometry technique."757" This well-established technique. described in ?.. allows esting of optical surfaces and lenses. and measurement of wave-front. deviations in the A/100 rawe,"," This well-established technique, described in \citet{1974ApOpt..13.2693B}, allows testing of optical surfaces and lenses, and measurement of wave-front deviations in the $\lambda/100$ range."758 Some systematic error sources such as extraueous [riuges. that is parasitic interference. were exaimiued iun the same paper auc more extensively in 2..," Some systematic error sources such as 'extraneous' fringes, that is parasitic interference, were examined in the same paper and more extensively in \citet{Schwider:83}."759 In our case. the theoretical description of this perturbation is hereafter detailed in the framework of a simple two-telescope," In our case, the theoretical description of this perturbation is hereafter detailed in the framework of a simple two-telescope"760 (Vokr«mublicky&Nesvoru*2008:Vokrouhllick*2009) (es.Binzel1996).. 13)," \citep{vok08,pra09} \citep[e.g.][]{bin96}. \ref{tbl-1})"761 27.500 (1 (seeTrilline2010.foranoverview). (Rabi id," $>7,500$ $\sim1$ \citep[see][for an overview]{tri10}, \citep{rab98}."762"s due to the lack of data of sub- NEAs,", due to the lack of data of sub-kilometer-size NEAs.763 Ou he other 1mand. the receutly ideutified ure-asteroid candidates are likely to boe of cohbunon origin (Pravec&Vokroihlicky2009).," On the other hand, the recently identified paired-asteroid candidates are likely to be of common origin \citep{pra09}."764. Tt coul be a convincing evidence to support this ivpothesis if both componeits within a pair are xovecd o have ideutical classificatio1. but however oivsiea observations are lacking Or almost all xuüred-asterokl eanlidates until now.," It could be a convincing evidence to support this hypothesis if both components within a pair are proved to have identical classification, but however physical observations are lacking for almost all paired-asteroid candidates until now."765" The moetho colbinine visual/near-intrarcd. spectroscopy and thermal iufrared measurement is preferred amone all practical eround-based nethods as it xovides hiehest accuracy as well as lost complete information of a target in nost cases, bith owever it d8 also very time CCunsunmiue alc eenerallv requires iuediuni or arge telescopes."," The method combining visual/near-infrared spectroscopy and thermal infrared measurement is preferred among all practical ground-based methods as it provides highest accuracy as well as most complete information of a target in most cases, but however it is also very time consuming and generally requires medium or large telescopes."766 By contrast. broad-band BVRT yhotomerv only allows crude classification. but it is more efücient than spectroscopy as it does not require as muuch times aud efforts as the latter. aud he resul can be useful for preliminary diagnose »urpose.," By contrast, broad-band $BVRI$ photometry only allows crude classification, but it is more efficient than spectroscopy as it does not require as much times and efforts as the latter, and the result can be useful for preliminary diagnose purpose."767 Iu this stidy. we emploved this method o investigate some selected. unusual asteroids.," In this study, we employed this method to investigate some selected unusual asteroids."768 Most of hese asteroids had not beeu reportedly classified prior to f16 observational pliase of this work., Most of these asteroids had not been reportedly classified prior to the observational phase of this work.769 Description «M the observation xocedure. data reduction aud details of classiHcatfion ds preseuted in Sectio 12 and 3.," Description of the observation procedure, data reduction and details of classification is presented in Section 2 and 3."770 We then conrpared and merged our restIts with other repored studies to assess the result consistenev. betwee1 onrs and others (Section [.1) and investigate the degrees of consistency with theoretical expectatiois (Section 1.2 ancl 1.3)., We then compared and merged our results with other reported studies to assess the result consistency between ours and others (Section 4.1) and investigate the degrees of consistency with theoretical expectations (Section 4.2 and 4.3).771 The Lulin Ouc-ineter Telescope (LOT) at Luliu Observatory. Tawa. was emploved for this study. except for one asteroid. 2008 EV. which wasobservedwith the0.[1-au telescopeof the same observatory.," The Lulin One-meter Telescope (LOT) at Lulin Observatory, Taiwan, was employed for this study, except for one asteroid, 2008 EV5, which wasobservedwith the0.41-m telescopeof the same observatory."772 The 0.[1-3 telescope observations for, The 0.41-m telescope observations for773]t has recently. become accepted. that quasar activity and black hole growth are an integral part of galaxy evolution. however a detailed: understanding of what triggers quasar activity and how they are fueled still clucles us.,"It has recently become accepted that quasar activity and black hole growth are an integral part of galaxy evolution, however a detailed understanding of what triggers quasar activity and how they are fueled still eludes us."774 The leading contender for the identity of luminous. high redshift QSOs is that they are black holes fed by by major mergers of gas- galaxies1990).," The leading contender for the identity of luminous, high redshift QSOs is that they are black holes fed by by major mergers of gas-rich galaxies."775. Recent incarnations of such mocels provide a good. description of many observed properties of the QSO population., Recent incarnations of such models provide a good description of many observed properties of the QSO population.776 The situation is particularly interesting at high redshift. where the population of supermassive black holes that powers the QSOs is growing rapidlyreview).," The situation is particularly interesting at high redshift, where the population of supermassive black holes that powers the QSOs is growing rapidly."777. To further understand this important phase of black hole and galaxy evolution we would like to build a model in which QSO activity is tied to the evolving cosmic web of dark matter halos., To further understand this important phase of black hole and galaxy evolution we would like to build a model in which QSO activity is tied to the evolving cosmic web of dark matter halos.778 The relationship between QSOs and dark matter halos. their environments and duty eveles. can be constrained. via observations of their space density ancl large-scale clustering2001).," The relationship between QSOs and dark matter halos, their environments and duty cycles, can be constrained via observations of their space density and large-scale clustering."779. These constraints become particularly sensitive if the QSOs inhabit the rarest. most massive halos for which the spatial clustering depends strongly on halo mass1989).," These constraints become particularly sensitive if the QSOs inhabit the rarest, most massive halos for which the spatial clustering depends strongly on halo mass."780. At redshifts 2<3 the advent of large optical surveys for QSOs has led to firm constraints on the clustering as a function of luminosity and. redshift2008)., At redshifts $z<3$ the advent of large optical surveys for QSOs has led to firm constraints on the clustering as a function of luminosity and redshift.781. With the Sloan Digital Sky Survey (SDSS) we are now able to measure the clustering of QSOs well even at 2>3mainDBodyCitationEnd242]She07., With the Sloan Digital Sky Survey (SDSS) we are now able to measure the clustering of QSOs well even at $z>3$.782 Interestingly. the correlation length of the QSO population increases rapidly with redshift. from ry=16.90d:1.73 at τς23 t0 ny=24.3042.36 at c4.," Interestingly, the correlation length of the QSO population increases rapidly with redshift, from $r_0=16.90\pm 1.73$ at $z\simeq 3$ to $r_0=24.30\pm 2.36$ at $z\simeq 4$."783 demonstrate that they can fit the observed. clustering and space density of 2~4 QSOs with the model of provided the z~4 QSOs are relatively long lived (Fo~160 Myr) and inhabit halos more massive than about 5.Loh!M..," demonstrate that they can fit the observed clustering and space density of $z\sim 4$ QSOs with the model of provided the $z\sim 4$ QSOs are relatively long lived $\tq\sim 160\,$ Myr) and inhabit halos more massive than about $5\times 10^{12}\,h^{-1}M_\odot$."784 For this caleulation assume that there is a monotonic relationship between instantaneous QSO luminosity and halo mass. with no scatter.," For this calculation assume that there is a monotonic relationship between instantaneous QSO luminosity and halo mass, with no scatter."785 Lhe actual relation between instantaneous. QSO luminosity ancl host halo mass is expected to include some scatter., The actual relation between instantaneous QSO luminosity and host halo mass is expected to include some scatter.786 Scatter is expected. in. several of the relationships linking the QSO luminosity ancl the host halo mass: the relationship between the host halo mass ancl galaxy. bulge. in the relationship between galaxy. bulge ancl black hole mass. in the relationship between black hole mass ancl peak luminosity and in the relationship between peak and instantaneous QSO Luminosity.," Scatter is expected in several of the relationships linking the QSO luminosity and the host halo mass: the relationship between the host halo mass and galaxy bulge, in the relationship between galaxy bulge and black hole mass, in the relationship between black hole mass and peak luminosity and in the relationship between peak and instantaneous QSO luminosity."787 llere we demonstrate that the very. high correlation length. measured. by(2007).. when combined with the rapid increase in bias for the most massive halos. stronely constrains the scatter between instantaneous QSO," Here we demonstrate that the very high correlation length measured by, when combined with the rapid increase in bias for the most massive halos, strongly constrains the scatter between instantaneous QSO"788where is the direction of the shear axis. aud j is a 1ueasure of shear.,"where $\beta$ is the direction of the shear axis, and $\eta$ is a measure of shear."789 The “couformal shear” j cau be reparameterized as the distortion ó=taulis or the reduced shear g=tanh5/2., The “conformal shear” $\eta$ can be reparameterized as the distortion $\delta=\tanh\eta$ or the reduced shear $g=\tanh\eta/2$.790 The shape must be assigned to au imuage of a ealaxv with some surface-brightuess distribution I(0)., The shape must be assigned to an image of a galaxy with some surface-brightness distribution $I(\boldtheta)$.791 Tnitially we will ignore the effects of PSF convolution ou the observed tuage., Initially we will ignore the effects of PSF convolution on the observed image.792" The BJ02 definition of shape is to specify roundness criteria or ""eieubluidtv tests.” AL(P) aud AL.(I). that operate on 7 to vield oue scalar for cach componcut of the slieartypically these are quadrupole moments."," The BJ02 definition of shape is to specify roundness criteria or “circularity tests,” $M_+(I)$ and $M_\times(I)$, that operate on $I$ to yield one scalar for each component of the shear—typically these are quadrupole moments."793" The object is deemed circular (e= 0) if AL,(f)=AL.(2)0.", The object is deemed circular $\bolde=0$ ) if $M_+(I)=M_\times(I)=0$.794 ΤΕ the object is not circular. then we assign to the object the shape wwhich vields the solutions we find the shear tthat. when applied to the coordinate svsteni@.. amakes the nage appear circular iu that coordinate system. and declare the galaxy shape to be this shear.," If the object is not circular, then we assign to the object the shape which yields the solutions we find the shear that, when applied to the coordinate system, makes the image appear circular in that coordinate system, and declare the galaxy shape to be this shear."795 αν circularity test will do. as long as has a unique solution13: iu particular the matrix ΝΤΑΙ iust be non-sugular.," Any circularity test will do, as long as has a unique solution; in particular the matrix $d{\bf M} / d\boldeta$ must be non-singular."796" The shape lis defined by ο=tanhy. keeping the positiou auele >,"," The shape is defined by $e=\tanh\eta$, keeping the position angle $\beta$."797" Defining shape in this wav with a suitable circularity test has the virtue that the effect of a lensing distortion ""upon the galaxy shape ls completely defined by the multiplication of shear matrices.", Defining shape in this way with a suitable circularity test has the virtue that the effect of a lensing distortion upon the galaxy shape is completely defined by the multiplication of shear matrices.798" In particular. the compoucnt-wise formulae for transformation of a shape under a shear must take the form given by Miralda-Escudé (1991):: We can take the limit of a weak shear à,1 à.Ξ0 DJO02 describe (855) a scheme for optimally weieliting and combining an eusemble of shapes to produce an accurate estimate of the distortion6.."," In particular, the component-wise formulae for transformation of a shape under a shear must take the form given by \citet{escude91}: We can take the limit of a weak shear $\delta_+\ll 1$, $\delta_\times =0$: BJ02 describe 5) a scheme for optimally weighting and combining an ensemble of shapes to produce an accurate estimate of the distortion."799 This scheme is predicated on the assignment of shapes that transform under shear according to(6)., This scheme is predicated on the assignment of shapes that transform under shear according to.800. IIeuce to test the accuracy of our uethodolosv im recovering weak lensing shear. Furthermore. he isotropy of the Universe euarautees that the e; of an uuleused population will be unuiforiulv distributed im. hence we need only verity hat Equations (7)) hold when averaged over au ensemble of galaxies with fixed unleused. fe] but random orientations: Tere the brackets refer to averaging over the pre-lensing oricutation J.," Hence to test the accuracy of our methodology in recovering weak lensing shear, Furthermore, the isotropy of the Universe guarantees that the $\bolde_i$ of an unlensed population will be uniformly distributed in $\beta$, hence we need only verify that Equations \ref{weakmap}) ) hold when averaged over an ensemble of galaxies with fixed unlensed $|e|$ but random orientations: Here the brackets refer to averaging over the pre-lensing orientation $\beta$."801" The secoud order term in Ó, vanishes. so these equations are valid to O(6?)."," The second order term in $\delta_+$ vanishes, so these equations are valid to $O(\delta^3)$."802 We refer to this as the (Fig. 1)).," We refer to this as the (Fig. \ref{fig:ringtest}) ),"803 since we coustruct an ensemble of test galaxics which form a rug in the pplaue. then shear them. measure their shapes. and take the mean.," since we construct an ensemble of test galaxies which form a ring in the plane, then shear them, measure their shapes, and take the mean."804 We also note that as a special case. we should obtain (e=0 when there is applied shear.," We also note that as a special case, we should obtain $\langle805\bolde^\prime\rangle=0$ when there is applied shear."806 When there is no PSF or the PSF Is svuunetric under rrotation. then this result holds for any measurement scheme that is sviunietric under inversion or exchange of the .c and y axes of inages.," When there is no PSF or the PSF is symmetric under rotation, then this result holds for any measurement scheme that is symmetric under inversion or exchange of the $x$ and $y$ axes of images."807 But for an asvuuuetrie PSF. this is a stringent test of the ability of the shape-measurement techuique to remove the effects of the PSF from the galaxy shapes.," But for an asymmetric PSF, this is a stringent test of the ability of the shape-measurement technique to remove the effects of the PSF from the galaxy shapes."808 The circularity test described in BJO2 involves decomposing the pre-couvolution surface brightucss distribution of the galaxw. 16]. into the Causs-Laguerre set of orthonormal basis functious iu the plane.," The circularity test described in BJ02 involves decomposing the pre-convolution surface brightness distribution of the galaxy, $I(\boldtheta)$, into the Gauss-Laguerre set of orthonormal basis functions in the plane."809 We consider here a general set of two dimensional functions {¢;(@)} that are complete (though not necessarily orthogonal) over the plane., We consider here a general set of two dimensional functions $\{\psi_i(\boldtheta)\}$ that are complete (though not necessarily orthogonal) over the plane.810 Auv set of complete fuuctious can be transformed to a new complete set for} via, Any set of complete functions can be transformed to a new complete set $\{\psi^E_i\}$ via811w Belezvisski and Bulik (1999).,by Belczyńsski and Bulik \shortcite{BB1998}.812. We concentrate on the dependence of the properties of the compact object binaries on the parameters used in the population synthesis code., We concentrate on the dependence of the properties of the compact object binaries on the parameters used in the population synthesis code.813 We ind that the most important parameter that determines the »opulation of compact object binaries is the kick velocity a neutron star receives at birth. however this distribution is poorly known.," We find that the most important parameter that determines the population of compact object binaries is the kick velocity a neutron star receives at birth, however this distribution is poorly known."814 [ρου and Tutukov (1996)— claim that he properties of pulsars can be explained by only the recoil velocities with no need for the kicks., Iben and Tutukov \shortcite{1996ApJ...456..738I} claim that the properties of pulsars can be explained by only the recoil velocities with no need for the kicks.815 Blaauw and tamachancran (1998) find that a single kick velocity. of 200 sullices to reproduce the pulsar population.," Blaauw and Ramachandran \shortcite{BlauwRama1998} find that a single kick velocity of $200\,$ $^{-1}$ suffices to reproduce the pulsar population."816 Cordes and Chernol! (1997). proposed a weighted sum two Gaussians: SO percent with the width 175 kk s and 20 percent with the width 700 km , Cordes and Chernoff \shortcite{1997ApJ...482..971C} proposed a weighted sum two Gaussians: 80 percent with the width $175$ km $^{-1}$ and 20 percent with the width $700$ km $^{-1}$.817We outline the model of the binary evolution. and propagation in a galactic potential in section 2., We outline the model of the binary evolution and propagation in a galactic potential in section 2.818 The results of the caleulation are presented in section 3 and we discuss them in section 4., The results of the calculation are presented in section 3 and we discuss them in section 4.819 In order to study the spatial distribution of compact object nmiergers we use the population svnthesis method., In order to study the spatial distribution of compact object mergers we use the population synthesis method.820 We use he population svnthesis code (Belezvnskiand.Bulik.1999) which concentrates on the population of massive star unaries. Le. those that may eventually lead: to formation of compact objects and compact object binaries.," We use the population synthesis code \cite{BB1998} which concentrates on the population of massive star binaries, i.e. those that may eventually lead to formation of compact objects and compact object binaries."821 We include he evolution of the binaries due to interaction and. mass ransfer and also the kicks that a newly born neutron star receives in supernova explosion., We include the evolution of the binaries due to interaction and mass transfer and also the kicks that a newly born neutron star receives in supernova explosion.822 A binary may be disrupted. in each of the supernova events., A binary may be disrupted in each of the supernova events.823 The surviving xnaries obtain center of mass velocities. which change their rajectories and may even eject them from their galaxy.," The surviving binaries obtain center of mass velocities, which change their trajectories and may even eject them from their galaxy."824 While the evolution. of single stars depends. only on heir mass and metallicity the evolution of binaries is also a 'unction the initial orbit (semimajor axis ev. and eccentricity c) of the two stars.," While the evolution of single stars depends only on their mass and metallicity the evolution of binaries is also a function the initial orbit (semimajor axis $a$, and eccentricity $e$ ) of the two stars."825 We assume that the distribution. of he initial parameters can be expressed as a product. of distributions of four parameters: the larger star (primary) mass AZ. the mass ratio of the less massive to the more massive star in the binary d. and the orbital paranictors α and ce. be that this quantities are independent.," We assume that the distribution of the initial parameters can be expressed as a product of distributions of four parameters: the larger star (primary) mass $M$, the mass ratio of the less massive to the more massive star in the binary $q$, and the orbital parameters $a$ and $e$, i.e that this quantities are independent."826 The distribution of primary masses used here is (Bethe3rown.1998) and we adopt a Hat distribution of the mass ratio q., The distribution of primary masses used here is \cite{Bethe1998} and we adopt a flat distribution of the mass ratio $q$.827 The semi major axis clistribution is scale invariant. i.e. with the limits GAR.<à6000/2.. and we craw the eccentricity [rom a distribution (e)=2e.," The semi major axis distribution is scale invariant, i.e. with the limits $6R_\odot < a < 6000R_\odot$, and we draw the eccentricity from a distribution $\Xi(e) = 2e$."828" We assume that the kick velocity distribution is a three dimensional Gaussian. and. parameterize it with its width σε. Le. We eenerate population of compact object binaries for a few values of o, in order to asses the sensitivity of our results to this parameter."," We assume that the kick velocity distribution is a three dimensional Gaussian, and parameterize it with its width $\sigma_v$, i.e. We generate population of compact object binaries for a few values of $\sigma_v$ in order to asses the sensitivity of our results to this parameter."829 We describe the mass transfer in the common envelope evolution by the common envelope parameter ace (see og. (Vranckenetab. 1991))). and we use an intermediate value of O.S for this parameter.," We describe the mass transfer in the common envelope evolution by the common envelope parameter $\alpha_{CE}$ (see e.g. \cite{1991A&A...249..411V}) ), and we use an intermediate value of $0.8$ for this parameter."830 In this type of evolution he more massive star looses its envelope ancl becomes a relium star with mass approximately of its initial value., In this type of evolution the more massive star looses its envelope and becomes a helium star with mass approximately of its initial value.831 The 3 parameter which describes the specific angular momentum of the material expelled from the binary in the toche lobe overllow phase is set to 3=6 (PolsandAlar-inus. 1994)., The $\beta$ parameter which describes the specific angular momentum of the material expelled from the binary in the Roche lobe overflow phase is set to $\beta=6$ \cite{1994A&A...288..475P}.832. Accretion onto a neutron star in a binary is reated as Boncli-Llovle accretion and we use the formalism developed by (BetheancBrown.1998). to find the amount of mass accreted onto the neutron star. and the final orbital separation.," Accretion onto a neutron star in a binary is treated as Bondi-Hoyle accretion and we use the formalism developed by \cite{Bethe1998}833 to find the amount of mass accreted onto the neutron star, and the final orbital separation."834 Systems with nearly equal masses evolve at the similar speed. and loose the common envelope. shrinking their orbit at the same time.," Systems with nearly equal masses evolve at the similar speed, and loose the common envelope, shrinking their orbit at the same time."835 For a more detailed description of the population svnthesis code see (Belezvnskiancl 19001., For a more detailed description of the population synthesis code see \cite{BB1998}.836 We assume that a neutron star with mass of 1.4. is formed in each supernova explosion.," We assume that a neutron star with mass of $1.4\, M_\odot$ is formed in each supernova explosion."837 We clraw a random time in the orbital motion to obtain the position on the orbit when the supernova explodes., We draw a random time in the orbital motion to obtain the position on the orbit when the supernova explodes.838 The remaining mass of the envelope is ejected from the svstem. and the newly formed neutron star receives a Kick.," The remaining mass of the envelope is ejected from the system, and the newly formed neutron star receives a kick."839Me verily whether the svstenm is still bound after the explosion.,We verify whether the system is still bound after the explosion.840" For bound systems we find the parameters of the new orbit and the kick velocity the whole binary receives where AZ, is the mass of the companion. ALS.ALM are the initial ancl final masses of the supernova. οὖν is the orbital velocity ofthe supernova at the time of explosion."," For bound systems we find the parameters of the new orbit and the kick velocity the whole binary receives where $M_1$ is the mass of the companion, $M_2^i,\,841M_2^f$ are the initial and final masses of the supernova, $\vec842v_2$ is the orbital velocity of the supernova at the time of explosion."843 After cach supernova expolsion we verily whther the svstem survives as a binary., After each supernova expolsion we verify whther the system survives as a binary.844 A compact object. binary loses its encrey through >eravitational radiation., A compact object binary loses its energy through gravitational radiation.845 Phe time to mergeὃν is (Peters. where e& is the semi major axis of the orbit. ¢ is its," The time to merge is \cite{Peters1964}846 where $a$ is the semi major axis of the orbit, $e$ is its"847second eroup) is less clear because several ecceutric ECPs with distances of closest as πα as 0.05 AU are also known.,second group) is less clear because several eccentric EGPs with distances of closest as small as $0.05$ AU are also known.848 It is thus uot clear why some ECPs with periastron distances larger than this value would be tidally-circularized while others would not., It is thus not clear why some EGPs with periastron distances larger than this value would be tidally-circularized while others would not.849 We note. however. that for values of the tidal parameter Q not too differcut from that of Jupiter (~10°). EGPs in this second group are also expected to be svuchrouized.," We note, however, that for values of the tidal parameter $Q$ not too different from that of Jupiter $\sim 10^5$ ), EGPs in this second group are also expected to be synchronized."850 We will asstuue it is indeed the case in our calculations., We will assume it is indeed the case in our calculations.851" Racial velocity surveys only measure Af,sin. which is a lower Inuit to the plauct’s mass. AZ. given the nuknown orbital inclination. 4."," Radial velocity surveys only measure $M_p \sin i$, which is a lower limit to the planet's mass, $M_p$, given the unknown orbital inclination, $i$."852" For randomly oriented svstenis. the distribution of when estimating the rage of likely values for R,, aud D, For a given mass. M,. the radius of an isolated planct eravitational acceleration is derived as g—GM, where G is the eravitational coustant."," For randomly oriented systems, the distribution of when estimating the range of likely values for $R_o$ and $B_u$ For a given mass, $M_p$, the radius of an isolated planet gravitational acceleration is derived as $g=GM_p/R_p^2$, where $G$ is the gravitational constant."853" For the m.mean laver thickness. f7. we adopt the atmospheric pressure scale-height. Mam,=ωνο. where Rois the perfect eas constaut."," For the mean layer thickness, $H$, we adopt the atmospheric pressure scale-height, $H_{\rm atm} \equiv {{\cal R} T_{\rm atm}}/{g}$, where ${\cal R}$ is the perfect gas constant."854" The elobal radiative equilibrium temperature of the planet is Dara,= which is a faction of the pareut star ΠΟ (CL,x T1R2). the planet που. a. and Bond albedo jy."," The global radiative equilibrium temperature of the planet is $T_{\rm atm} = T_\star855({R_\star}/{2 a})^{1/2} (1-A_b)^{1/4}$ , which is a function of the parent star luminosity $L_\star \propto T_\star^4 R_\star^2$ ), the planet's semi-major axis $a$, and Bond albedo $A_b$."856" We adopt A,=0.5 for all our ποΊσα estimates: our results only weakly depend onhe value of Ay unless it approaches unity.", We adopt $A_b=0.5$ for all our numerical estimates; our results only weakly depend onthe value of $A_b$ unless it approaches unity.857" The stellar tuinosity is derived from the mass through the simple nass-huninosity relation £,=(AL/ALy?9L...", The stellar luminosity is derived from the mass through the simple mass-luminosity relation $L_\star =(M_\star / M_\odot)^{3.6} L_\odot$.858" The last two parameters needed to determine &, aud D, ave the planetary rotation rate O and the global kinetic ΟΠΟΙΟΥ scale ££.", The last two parameters needed to determine $R_o$ and $B_u$ are the planetary rotation rate $\Omega$ and the global kinetic energy scale $\bar U$.859" We asstune that Q=O4, (as determined o» radial velocity surveys) in all cases.", We assume that $\Omega =\Omega_{\rm orb} $ (as determined by radial velocity surveys) in all cases.860 We allow C to vary from 50 urs +. the suallest observed value for giaut auets iu the Solar System (Jupiter). to 1000 12 lon rather large value for which the tvpical wind speeds in the atmosphere of hot. close-in EGPs approaches the souud speed.," We allow $\bar U$ to vary from $50$ m $^{-1}$, the smallest observed value for giant planets in the Solar System (Jupiter), to $1000$ m $^{-1}$ , a rather large value for which the typical wind speeds in the atmosphere of hot, close-in EGPs approaches the sound speed."861" A value C=400 31 + is adopted for our fiducial estimate of R, aud D,.", A value $\bar U =400$ m $^{-1}$ is adopted for our fiducial estimate of $R_o$ and $B_u$.862" Estimated values for 2, aud B, are given in Table 1 for Solar System giauts aud close-in ECPs."," Estimated values for $R_o$ and $B_u$ are given in Table \ref{tab:one}863 for Solar System giants and close-in EGPs."864 The values listed for close-in EGPs correspond to the range of mii./max., The values listed for close-in EGPs correspond to the range of min./max.865 values found elven the various asstuuptions detailed iu 833., values found given the various assumptions detailed in 3.866 Fiducial estimates are also reported iu figure 1.. where solid. dots correspond to group 1. ECGPs (sate tidal svuchnrouization assuniptiou) and open circles to eroup 2 EGPs (tidal svuchronization assunptiou less sate).," Fiducial estimates are also reported in figure \ref{fig:one}, where solid dots correspond to group 1 EGPs (safe tidal synchronization assumption) and open circles to group 2 EGPs (tidal synchronization assumption less safe)."867 IID 209158 b is indicated as a star., HD 209458 b is indicated as a star.868" It is clear from figure 1. that closean EGPs. as a eroup. occupy a ciffercut region of the R,, B, parameter space than Solar System giants."," It is clear from figure \ref{fig:one} that close-in EGPs, as a group, occupy a different region of the $R_o$ $B_u$ parameter space than Solar System giants."869" In particular. thev systematically have a Bureer unuuber DB,7»1/9 (oven when accounting for the large range of allowed values: Table 1)). which indicates that the preseuce of eirciunpolar vortices is expected in the radiative region within the framework of shallow-water dvuanuies."," In particular, they systematically have a Burger number $B_u > 1/9$ (even when accounting for the large range of allowed values; Table \ref{tab:one}) ), which indicates that the presence of circumpolar vortices is expected in the radiative region within the framework of shallow-water dynamics."870" The lareer values of R, also indicate that generally few bauds/jets are expected ou these planets (the πιοταλατν ou C stronglv affects this nuuber: see Table 1)). thus allowing the formation of larger ""great spots (πάσι could also contribute to the variability: Cho et al."," The larger values of $R_o$ also indicate that generally few bands/jets are expected on these planets (the uncertainty on $\bar U$ strongly affects this number; see Table \ref{tab:one}) ), thus allowing the formation of larger “great spots” (which could also contribute to the variability; Cho et al."871 2002h)., 2002b).872" The near aliguinent of all the points represcuting close-in ECPs in figure 1. shows that the dominant parameter determining their position in this diagram is their rotation rate (R,x»οἘν D,xο7)."," The near alignment of all the points representing close-in EGPs in figure \ref{fig:one} shows that the dominant parameter determining their position in this diagram is their rotation rate $R_o \propto \Omega^{-1}$; $B_u \propto873\Omega^{-2}$ )."874" The simall values of F2, and D, for Solar Svstem eiauts reflect their relatively fast rotation rates.", The small values of $R_o$ and $B_u$ for Solar System giants reflect their relatively fast rotation rates.875 Although we argued in favor of variable atmospheric siguatures for close-in EGPs. it is important to note hat models do not vet quantitatively predict how much variability is expected.," Although we argued in favor of variable atmospheric signatures for close-in EGPs, it is important to note that models do not yet quantitatively predict how much variability is expected."876 Iun Cho et al. (, In Cho et al. (8772002a.b). we enmlhiasized that the combination of C (unknown) and he amplitude of dav-uight heating (parametrized iu aciabatic simulatious) determines the contrast of the hermal spots associated with circmmpolar vortices.,"2002a,b), we emphasized that the combination of $\bar U$ (unknown) and the amplitude of day-night heating (parametrized in adiabatic simulations) determines the contrast of the thermal spots associated with circumpolar vortices."878 In the ture. diabatie shallow-water models will allow a scl determination of the dav-night forcing.," In the future, diabatic shallow-water models will allow a self-consistent determination of the day-night forcing."879 to make quantitative predictions regarding the level of variability expected for various atmosphericsignatures., to make quantitative predictions regarding the level of variability expected for various atmosphericsignatures.880 Support for this work was provided by NASA through Chandra Fellowship grant PFE9-10006 awarded by the Sinithsonian Astroplivsical Observatory for NASA uuder contract NASS-39073., Support for this work was provided by NASA through Chandra Fellowship grant PF9-10006 awarded by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-39073.881spectra (by definition). it is possible that the galaxies are heavily clust-extineted (Shiova&Bekki2000).. or that weak residual star formation activity exists associated with the detected radio emission. which is not strong enough to produce optical emission lines.,"spectra (by definition), it is possible that the galaxies are heavily dust-extincted \citep{shioya00}, or that weak residual star formation activity exists associated with the detected radio emission, which is not strong enough to produce optical emission lines."882" The timescale of racio emission from star formation can range from about 10"" to LO vr. depending on the supply and. escape of cosmic ravs and their environment as well as the lifetime of LLL regions (Chi&Wollendale1990:HelouBieay1993)."," The timescale of radio emission from star formation can range from about $10^{7}$ to $10^{8}$ yr, depending on the supply and escape of cosmic rays and their environment as well as the lifetime of HII regions \citep{chi90,helou93}."883. This timescale might be lone enough to be an explanation, This timescale might be long enough to be an explanation884conibined these suberoups under their main eroup type and come up with a total of 11 major categories.,combined these subgroups under their main group type and come up with a total of 11 major categories.885 We have assigned categories for all matched sources aud include the source types in Tables 2.0L., We have assigned categories for all matched sources and include the source types in Tables 2 – 4.886 We note that (228) of the stars are not assigned candidate types. either because they fall between the chosen category boundaries (true for 160 sources) or are saturated iu one or more 2ATASS buds.," We note that (228) of the stars are not assigned candidate types, either because they fall between the chosen category boundaries (true for 160 sources) or are saturated in one or more 2MASS bands."887 Six-duueusional data is difficult το represent: additionally. it is uot clear that all of the colors vield completely independent information.," Six-dimensional data is difficult to represent; additionally, it is not clear that all of the colors yield completely independent information."888 To simplify the analysis we have chosen the three axes which mest clearly illustrate the object type distinctions., To simplify the analysis we have chosen the three axes which most clearly illustrate the object type distinctions.889 Figure Bois alAu. We A color-color diagram. where the colors of the poiuts represents the /7A. color.," Figure 3 is a $J-K_{s}$ $K_{s}-$ A color-color diagram, where the colors of the points represents the $H-K_{s}$ color."890 Photometric uncertainties (riu) in cach color are also shown., Photometric uncertainties (rms) in each color are also shown.891" The JIv, NM, A color plane shows several obvious groupings. while the JFA, color is able to wither separate the other groups."," The $J-K_{s}$ $K_{s}-$ A color plane shows several obvious groupings, while the $H-K_{s}$ color is able to further separate the other groups."892 To identity the sources contained iu the cluster-analvsis categories. woe lave exiunined the colors expected for the various source types w the Wainscoat et al. (," To identify the sources contained in the cluster-analysis categories, we have examined the colors expected for the various source types by the Wainscoat et al. ("893"1992) ""SIKY"" model of the Calaxy and compared these to the observed colors.",1992) “SKY” model of the Galaxy and compared these to the observed colors.894 Absolute naenitudes in the MSN A baud were supplied for cach of he 87 objects by M. Cohen (priv., Absolute magnitudes in the MSX A band were supplied for each of the 87 objects by M. Cohen (priv.895 comun.)., comm.).896 FFigure LE is he model analog to Figure 3., Figure 4 is the model analog to Figure 3.897 The 1l categories occupyiug the threc-axis color-color diagram in Figure 3 and the number of sources assigucd ο cach eroup. are as follows: Category V contains a mixture of several types of stars.," The 11 categories occupying the three-axis color-color diagram in Figure 3 and the number of sources assigned to each group, are as follows: Category V contains a mixture of several types of stars."898 From Wainscoat et al. (, From Wainscoat et al. (8991992) we see that the colors of red supergiauts (RSCs) and carly oxveeu-rich aud carbou-rich ACB stars are all quite simibku. (,1992) we see that the colors of red supergiants (RSGs) and early oxygen-rich and carbon-rich AGB stars are all quite similar. (900"By ""earlv we mean those stars with relatively little mass loss or obscuration of the central star at visible and ucar-IR waveleuetlis).",By “early” we mean those stars with relatively little mass loss or obscuration of the central star at visible and near-IR wavelengths).901 ffA.Toalesserdegree.wethenusethe color to distinguish between these objects the RSCs being the bluest objects imong this group on this axis. aud the carbon stars beime the reddest.," To a lesser degree, we then use the $%902H-K_{s color to distinguish between these objects, the RSGs being the bluest objects among this group on this axis, and the carbon stars being the reddest."903 However. we do uot claim that these desigunatious are exact.," However, we do not claim that these designations are exact."904 In fact. even using all six color axes. some confusion still exists between some of these sources.," In fact, even using all six color axes, some confusion still exists between some of these sources."905 The designations below try to iiake the best separation between subtypes in this category: Examining just the JoWy Wy A two-color ciagrai. category VI appears to be a monolithic group.," The designations below try to make the best separation between subtypes in this category: Examining just the $J-K_{s}$ $K_{s}-$ A two-color diagram, category VI appears to be a monolithic group."906" The six-dimensional analysis of the observed colors aud the SIVY model reveal that it actually contains two cistiuct types of objects: the “late” (that is. high mass-loss rate. high obscuration) C-rich AGB stars (also known as ""infrared"" carbon stars: Chan νοκ 1988) aud their O-rich counterparts (generally OTL/TR stars and IIO maser sources)."," The six-dimensional analysis of the observed colors and the SKY model reveal that it actually contains two distinct types of objects: the “late” (that is, high mass-loss rate, high obscuration) C-rich AGB stars (also known as “infrared” carbon stars; Chan Kwok 1988) and their O-rich counterparts (generally OH/IR stars and $_{2}$ O maser sources)."907 Iu this category. the mixing is not as thorough as in category V. so the Corich aud O-vich sources are more casily separated.," In this category, the mixing is not as thorough as in category V, so the C-rich and O-rich sources are more easily separated."908 However. we still expect some confusion in the overlap region of these categories: Category VII represcuts objects with rather muusual colors.," However, we still expect some confusion in the overlap region of these categories: Category VII represents objects with rather unusual colors."909 The near-IR colors are blue. aud ou the JiIv. color-color diagram they die amoue the normal stars.," The near-IR colors are blue, and on the $JHK_{s}$ color-color diagram they lie among the normal stars."910 Iowever. they show a very large excess clissiou component at 8.3 422. indicative of dust associated with a stellar point source.," However, they show a very large excess emission component at 8.3 $\mu $ m, indicative of dust associated with a stellar point source."911 The SKY iuodel colors aud the SIMBAD ideutificatious show that the associatious of the 2ATASS sources and the MSX sources are indeed real., The SKY model colors and the SIMBAD identifications show that the associations of the 2MASS sources and the MSX sources are indeed real.912 The objects in this eroup teud to be either planctary nebulae (PNe). or carly (O- and B-type) stars associated with dust (1.0... IIT regions. reflection nebulae. LDVs. dusty D stars. ete).," The objects in this group tend to be either planetary nebulae (PNe), or early (O- and B-type) stars associated with dust (i.e., HII regions, reflection nebulae, LBVs, dusty B stars, etc.)."913" To some degree we can separate the objects based ou the ΠA, color. with the PNe being redder than the other objects."," To some degree we can separate the objects based on the $H-K_{s}$ color, with the PNe being redder than the other objects."914 However. these objects tend to be among the fainterof the 2\LASS sources.," However, these objects tend to be among the fainterof the 2MASS sources."915" Therefore. confusion still exists in ideutifving these sources ucar the assumed TPK, bounclary:"," Therefore, confusion still exists in identifying these sources near the assumed $H-K_{s}$ boundary:"916aanalysis compared to the aanalysis.,analysis compared to the analysis.917 QSOs below the SNR threshold are retained only as central QSOs (see Figure 49) in the correlation analysis. i.e.. the sightlines provided by their spectra are not included in the absorption system search or in the control absorber counts.," QSOs below the SNR threshold are retained only as central QSOs (see Figure \ref{fig:method}) ) in the correlation analysis, i.e., the sightlines provided by their spectra are not included in the absorption system search or in the control absorber counts."918 Next. we select the sightlines which we will use in the analysis.," Next, we select the sightlines which we will use in the analysis."919 We only consider sightlines from QSOs above the SNR threshold., We only consider sightlines from QSOs above the SNR threshold.920 As indicated in Figure + we restrict the sightlines to start at As>0.1 from their QSO host. and to lie redwards of iin the QSO rest frame tto avoid the strong Lyman-a line at aand the Lyman-a forest bluewards of Lyman-a).," As indicated in Figure \ref{fig:method} we restrict the sightlines to start at $\Dz>0.1$ from their QSO host, and to lie redwards of in the QSO rest frame (to avoid the strong $\alpha$ line at and the $\alpha$ forest bluewards of $\alpha$ )."921 The velocity difference criterion A>>0.1 ensures that the final signal is not affected by the excess of associated absorbers with sub-relativistic velocities described in Section ??.., The velocity difference criterion $\Dz>0.1$ ensures that the final signal is not affected by the excess of associated absorbers with sub-relativistic velocities described in Section \ref{sec:los}.922 Because of the restriction that sightlines must end at wwe cannot increase this limit much further without drastically reducing the number of QSO- aabsorber pairs., Because of the restriction that sightlines must end at we cannot increase this limit much further without drastically reducing the number of QSO- absorber pairs.923 Contained within these sightlines are 3374 aand 13504 aabsorbers., Contained within these sightlines are 3374 and 13504 absorbers.924 Using these sightlines and absorbers. we calculate the number density of absorbers as a function of redshift tz) shown in Figure 5.. and use this to estimate the expected number of absorbers in each A> bin along each sightline.," Using these sightlines and absorbers, we calculate the number density of absorbers as a function of redshift $n$ $z$ ) shown in Figure \ref{fig:nz}, and use this to estimate the expected number of absorbers in each $\Delta z$ bin along each sightline."925 Then. for each central QSO we use the sightlines of all its paired QSOs to measure the expected and observed number of absorbers as a function of comoving separation between the central QSO and the position along thes," Then, for each central QSO we use the sightlines of all its paired QSOs to measure the expected and observed number of absorbers as a function of comoving separation between the central QSO and the position along the."926ightline?.. In Figure 4. the arrow. labelled 7. indicates the comoving distance between the central QSO and one of the bins along one of the sightlines.," In Figure \ref{fig:method} the arrow, labelled $r$, indicates the comoving distance between the central QSO and one of the bins along one of the sightlines."927 In this bin. no absorber is found so no addition is made to the observed number of absorbers at this comoving separation.," In this bin, no absorber is found so no addition is made to the observed number of absorbers at this comoving separation."928 However. a small addition will be made to the expected number of absorbers at this separation. based on the pre-calculated Πές) at the redshift of this A> bin.," However, a small addition will be made to the expected number of absorbers at this separation, based on the pre-calculated $n$ $z$ ) at the redshift of this $\Delta z$ bin."929 As well as accruing the observed and expected numbers of absorbers as a function of comoving separation. we sum the number of contributing sight lines. which is necessary in the estimation of the errors.," As well as accruing the observed and expected numbers of absorbers as a function of comoving separation, we sum the number of contributing sight lines, which is necessary in the estimation of the errors."930 Finally. because we must bin our results in Av. and we do not have enough sightlines to create infinitesimal bins. we must account for the fact that more sightlines within a Ar range will lie towards the outer edge of the range than the inner edge i.e. the effective radius of our Ar bin is not centered on to the bin. but biased towards the outer edge.," Finally, because we must bin our results in $\Delta r$, and we do not have enough sightlines to create infinitesimal bins, we must account for the fact that more sightlines within a $\Delta r$ range will lie towards the outer edge of the range than the inner edge i.e. the effective radius of our $\Delta r$ bin is not centered on to the bin, but biased towards the outer edge."931 We do this by accumulating the expected number of absorbers. weighted by rLU (see Equation 3)).," We do this by accumulating the expected number of absorbers, weighted by $r^{-1.7}$ (see Equation \ref{eq:reff}) )."932 The final 3-D correlation estimate is given by: where (ο) is the observed number of absorbers in a Ay bin with effective radius raj. and Ni(r4) is the expected number as described above.," The final 3-D correlation estimate is given by: where ${\rm N_o(r_{eff})}$ is the observed number of absorbers in a $\Delta r$ bin with effective radius ${\rm r_{eff}}$, and ${\rm933N_e(r_{eff})}$ is the expected number as described above."934 The effective radius of each of the Ar bins is estimated to be: where the sum is over all As bins (segments of the sightlines in Figure 4) that contribute to the Ar bin., The effective radius of each of the $\Delta r$ bins is estimated to be: where the sum is over all $\Delta z$ bins (segments of the sightlines in Figure \ref{fig:method}) ) that contribute to the $\Delta r$ bin.935 + is chosen to be 1.7. which is close to the measured exponent in the 11--QSO correlation (see below). although the precise value used does not change the results significantly.," $\gamma$ is chosen to be 1.7, which is close to the measured exponent in the -QSO correlation (see below), although the precise value used does not change the results significantly."936 On small scales the QSO-absorber pairs in each Ay bin are independent (in general. absorbers are only paired to one QSO on small scales). thus the error is estimated assuming Poisson Statistics: On larger scales the pairs in each Ar bin are no longer independent (one absorber can be paired with many QSOs in the same bin) and the Poisson errors underestimate the true errors.," On small scales the QSO-absorber pairs in each $\Delta r$ bin are independent (in general, absorbers are only paired to one QSO on small scales), thus the error is estimated assuming Poisson statistics: On larger scales the pairs in each $\Delta r$ bin are no longer independent (one absorber can be paired with many QSOs in the same bin) and the Poisson errors underestimate the true errors."937" Following Shanks&Boyle¢1994).. as the number of pairs approaches the number of absorbers in the analysis (N,i.). the errors are approximated by: Finally. we calculate the significance of a positive clustering Signal (excess number of pairs) assuming a Poisson distribution,"," Following \citet{1994MNRAS.271..753S}, as the number of pairs approaches the number of absorbers in the analysis $_{\rm abs}$ ), the errors are approximated by: Finally, we calculate the significance of a positive clustering signal (excess number of pairs) assuming a Poisson distribution."938" Assuming no clustering. the probability of observing Ν,, or more pairs at any given r is given by: In the closest bins with fewest observed pairs. we checked that the probabilties derived here agree with those derived from a Binomial distribution."," Assuming no clustering, the probability of observing $_{\rm o}$ or more pairs at any given $r$ is given by: In the closest bins with fewest observed pairs, we checked that the probabilties derived here agree with those derived from a Binomial distribution."939 The correlation functions for aand aabsorber-QSO pairs are shown in Figure 6.. where we plot (1| £) in bins of ~ GMMpc.," The correlation functions for and absorber-QSO pairs are shown in Figure \ref{fig:trans}, where we plot $1+\xi$ ) in bins of $\sim6$ Mpc."940 The results are plotted logarithmically in the right hand panels., The results are plotted logarithmically in the right hand panels.941 As expected. there is an enhancement in the counts of both ppairs and ppairs at small comoving distances from the QSO.," As expected, there is an enhancement in the counts of both pairs and pairs at small comoving distances from the QSO."942 The significance of the detection is considerably higher for the aabsorbers. because the sample is much larger.," The significance of the detection is considerably higher for the absorbers, because the sample is much larger."943 Overplotted as a dashed line is the confidence limit for detection of a clustering signal. given the observed and expected number of absorbers in each bin (Eq. 63).," Overplotted as a dashed line is the confidence limit for detection of a clustering signal, given the observed and expected number of absorbers in each bin (Eq. \ref{eq:binomial}) )."944 In other words. if there were no clustering. we would expect an independent bin to lie above this line of the time.," In other words, if there were no clustering, we would expect an independent bin to lie above this line of the time."945 We can evaluate the amplitude and significance of the clustering signal for the, We can evaluate the amplitude and significance of the clustering signal for the946ciwarl galaxy masses (10175TALL) to massive galaxy clusters (1077Pf1 M.) across a large redshift range (2=0 ο 5).,dwarf galaxy masses $10^{10}\hmsol$ ) to massive galaxy clusters $10^{15}\hmsol$ ) across a large redshift range $z=0$ to $5$ ).947 We explored the orbital parameters of satellite halos at infall. when they cross within the virial racius of a larger vost halo.," We explored the orbital parameters of satellite halos at infall, when they cross within the virial radius of a larger host halo."948 The main results are as follows: As explored in refsee:cistribution.. the orbital distributions we find at z=0 when stacking halos of all masses agree well with those of previous work.," The main results are as follows: As explored in \\ref{sec:distribution}, the orbital distributions we find at $z=0$ when stacking halos of all masses agree well with those of previous work."949 Llowever. previous results of possible mass and/or redshift dependence of satellite orbits are mixed.," However, previous results of possible mass and/or redshift dependence of satellite orbits are mixed."950 ? founcl evidence that satellite orbits become more radial at higher halo mass scales but was unable to quantify his further. while ? and ? saw no such halo mass dependence over a limited mass range.," \citet{Ben05} found evidence that satellite orbits become more radial at higher halo mass scales but was unable to quantify this further, while \citet{VitKlyKra02} and \citet{WanJinMao05} saw no such halo mass dependence over a limited mass range."951 By contrast. ? found hat satellite angular momentum decreases with increasing satellite mass to host halo mass ratio. while 7? and 7? ound no dependence on mass ratio. though again over a imited mass ratio range.," By contrast, \citet{VitKlyKra02} found that satellite angular momentum decreases with increasing satellite mass to host halo mass ratio, while \citet{WanJinMao05} and \citet{KhoBur06} found no dependence on mass ratio, though again over a limited mass ratio range."952 Finally. the results of 2 suggested rends with redshift while those o£ 2. cid not.," Finally, the results of \citet{Ben05} suggested trends with redshift while those of \citet{VitKlyKra02} did not."953" La most cases. oevious work was limited in terms of merger statistics (in some cases. examining infall into a handful of halos) and dynamical range (unable to explore both Al5.Al, and Al«x A)."," In most cases, previous work was limited in terms of merger statistics (in some cases, examining infall into a handful of halos) and dynamical range (unable to explore both $M \gg M_*$ and $M \ll M_*$ )."954 The results here demonstrate clear dependence of satellite orbits on host halo mass ancl redshift’ and. no significant evidence for dependence on satellite mass., The results here demonstrate clear dependence of satellite orbits on host halo mass and redshift and no significant evidence for dependence on satellite mass.955 This is broadly consistent with the predictions of analytical triaxial collapse models (277).. in which more massive halos arise from a more spherical Lagrangian volume with less angular momentum.," This is broadly consistent with the predictions of analytical triaxial collapse models \citep{BBKS86,EisLoe95,SheMoTor01}, in which more massive halos arise from a more spherical Lagrangian volume with less angular momentum."956 We emphasize that the existenceof mass and redshift dependence implies that fits to orbital distributions based on stacking halos of all masses at ο~O (ο.2?) are not universally accurate.," We emphasize that the existenceof mass and redshift dependence implies that fits to orbital distributions based on stacking halos of all masses at $z\sim0$ \citep[e.g.,][]{Ben05,ZenBerBul05,WanJinMao05} are not universally accurate."957 The mass and redshift trends seen here have implications for various aspects of galaxy formation anc evolution., The mass and redshift trends seen here have implications for various aspects of galaxy formation and evolution.958 For example. recent work suggests that galaxy formation at 2 proceeds through narrow streams of colc eas (27).. fundamentally dillerent behavior than seen in the local Universe.," For example, recent work suggests that galaxy formation at $z \gtrsim 2$ proceeds through narrow streams of cold gas \citep{KerKatWei05,DekBirEng08}, fundamentally different behavior than seen in the local Universe."959 argued that this dillerence arises in par because massive 7.(1017 TAL.) halos at high redshilt form at the intersection of narrow filaments. while at low recdshif such halos are more likely embedded within a filament anc experience. wide-angle inflow.," \citet{DekBir06} argued that this difference arises in part because massive $\sim10^{12}\hmsol$ ) halos at high redshift form at the intersection of narrow filaments, while at low redshift such halos are more likely embedded within a filament and experience wide-angle inflow."960 Our results on the redshif evolution of satellite accretion qualitatively support this xeture. but it is not clear that the evolution seen. here is strong enough to imply a Fundamental change the nature of accretion at z2.5. when most orbital parameters diller rom their 50 values by ~204.," Our results on the redshift evolution of satellite accretion qualitatively support this picture, but it is not clear that the evolution seen here is strong enough to imply a fundamental change the nature of accretion at $z\sim2.5$, when most orbital parameters differ from their $z=0$ values by $\sim20\%$."961 This suggests. that he above results are driven more strongly by evolving gas hvsies than by the nature of mass accretion., This suggests that the above results are driven more strongly by evolving gas physics than by the nature of mass accretion.962 Our results also have clear implications for satellite galaxy evolution within eroups/clusters., Our results also have clear implications for satellite galaxy evolution within groups/clusters.963 For example. using he parametrization of satellite survival time given in 2.. changes in circularity across mass and redshift’ from the &lobal z=0 value can lead to reductions of satellite Lifetimes of 830% or more.," For example, using the parametrization of satellite survival time given in \citet{BoyMaQua08}, changes in circularity across mass and redshift from the global $z=0$ value can lead to reductions of satellite lifetimes of $30\%$ or more."964" Additionally. environmental ellects such as ram-pressture stripping of eas and tidal stripping of stars are expected to oceur primarily at orbital pericenter (e.g..2??).,"," Additionally, environmental effects such as ram-pressure stripping of gas and tidal stripping of stars are expected to occur primarily at orbital pericenter \citep[e.g.,][]{DekDevHet03,TayBab04,McCFreFon08}."965 3evond their dependence on evolving gas physies. our results sugeest that satellite galaxy quenching ancl morphological evolution proceed more cllicienthy ancl rapidly at higher eroup/cluster masses and higher redshift.," Beyond their dependence on evolving gas physics, our results suggest that satellite galaxy quenching and morphological evolution proceed more efficiently and rapidly at higher group/cluster masses and higher redshift."966 While this work focuses on the orbital parameters. of satellite halos at the time of infall. it is not immediately clear how well these orbital distributions and their mass ancl redshift’ dependencies persist to satellite populations well after infall as the orbits become allected by dynamical friction. tidal stripping. and extended: halo potentials (see the Appendix for some estimate of the latter).," While this work focuses on the orbital parameters of satellite halos at the time of infall, it is not immediately clear how well these orbital distributions and their mass and redshift dependencies persist to satellite populations well after infall, as the orbits become affected by dynamical friction, tidal stripping, and extended halo potentials (see the Appendix for some estimate of the latter)."967 There has been some work examining satellite subhalo orbits within host halos (???).. which found. orbital parameter distributions similar to those here (for example. 1)z0.5).," There has been some work examining satellite subhalo orbits within host halos \citep{GilKneGib04,ReeGovQui05,SalNavAba07}, which found orbital parameter distributions similar to those here (for example, $\eta\approx0.5$ )."968 This lack of circularity evolution after infall is plausible since dynamical friction is expected. to be inellicient. in altering orbital parameters such as circularity (2)..., This lack of circularity evolution after infall is plausible since dynamical friction is expected to be inefficient in altering orbital parameters such as circularity \citep{vdBLewLak99}.969 Furthermore. ? recently found that the orbits of satellite subhalos within host halos at z=0 are more racially biased in more massive host halos. which suggests that the results here remain valid well after infall.," Furthermore, \citet{Fal10} recently found that the orbits of satellite subhalos within host halos at $z=0$ are more radially biased in more massive host halos, which suggests that the results here remain valid well after infall."970 We will pursue a more robust investigation into the evolution of satellite orbits after infall in future work., We will pursue a more robust investigation into the evolution of satellite orbits after infall in future work.971 There is also possible observational evidence in support of the trends seen here., There is also possible observational evidence in support of the trends seen here.972 ? examined the orbital velocities of ealaxies in local galaxy. clusters. finding highly asymmetric velocity. distributions consistent with satellites [largely retaining their infalline orbits.," \citet{HerZarJin08} examined the orbital velocities of galaxies in local galaxy clusters, finding highly asymmetric velocity distributions consistent with satellites largely retaining their infalling orbits."973 Promisinglv. 2? examined satellite galaxy orbits in galaxy clusters from ο=0 to >=0.5 and found evidence that satellite orbits are indeed less isotropic (more racial) at higher redshilt.," Promisingly, \citet{BivPog09} examined satellite galaxy orbits in galaxy clusters from $z=0$ to $z=0.8$ and found evidence that satellite orbits are indeed less isotropic (more radial) at higher redshift."974 Finally. the results on satellite velocities also have implications for relating satellite clynamics to those of the overall eroup/cluster.," Finally, the results on satellite velocities also have implications for relating satellite dynamics to those of the overall group/cluster."975" At lower host halo masses. satellite velocities become significantly ""hotter. than than the host halo. implving a possible systematic biasing in using satellite velocity dispersions to infer halo masses (?.foundasimilar halos).."," At lower host halo masses, satellite velocities become significantly `hotter' than than the host halo, implying a possible systematic biasing in using satellite velocity dispersions to infer halo masses \citep[][found a similar trend with mass for the velocity bias of satellite subhalos within host halos]{Fal10}. ."976 Future work will involve a more detailed analysis of satellite velocity bias ancl its mass and. recishift dependence., Future work will involve a more detailed analysis of satellite velocity bias and its mass and redshift dependence.977 ] gratefully acknowledge the support of an NSE Graduate Research Fellowship., I gratefully acknowledge the support of an NSF Graduate Research Fellowship.978 E thank Martin. White for use of, I thank Martin White for use of979The data shown in Figure 4 show a clear trend for increasing dispersion with metal abundance.,The data shown in Figure 4 show a clear trend for increasing velocity dispersion with decreasing metal abundance.980 This is similar velocityto what is found by decreasingJohnsonetal.(2011) in their higher resolution data., This is similar to what is found by \cite{johnson10} in their higher resolution data.981" Babusiauxetal.(2010) uses two fields at b=—6? and b.=—12° and although there are few stars at [Fe/H] «—1 in their data, there is a hint of an increase in the velocity dispersion at lower metal abundances, in agreement with our observations."," \cite{babusiaux10} uses two fields at $b=-6^{\circ}$ and $b=982-12^{\circ}$ and although there are few stars at [Fe/H] $< -1$ in their data, there is a hint of an increase in the velocity dispersion at lower metal abundances, in agreement with our observations."983" Vieiraetal.(2007) find a flat distribution of metallicity with velocity dispersion (from proper motion data) for stars with [Fe/H] >—1, which is not in disagreement with our observations, where most of the increase in velocity dispersion takes place for lower stars. For metall"," \cite{vieira07} find a flat distribution of metallicity with velocity dispersion (from proper motion data) for stars with [Fe/H] $> -1$, which is not in disagreement with our observations, where most of the increase in velocity dispersion takes place for lower metallicity stars."984"icitystars with [Fe/H] >—1 the velocity dispersion is in good agreement with that measured for BRAVA M giants and does not depend strongly on metal abundance, which is broad with what measured Vieiraetal.(2007) in Plaut's agreementWindow and the two lower bygalactic latitude fields in Babusiauxetal.(2010)."," For stars with [Fe/H] $ > -1$ the velocity dispersion is in good agreement with that measured for BRAVA M giants and does not depend strongly on metal abundance, which is broad agreement with what measured by \cite{vieira07} in Plaut's Window and the two lower galactic latitude fields in \cite{babusiaux10}."985". However, at [Fe/H] «—1 there appears to be a (at face value) smooth transition to a dynamically hot component."," However, at [Fe/H] $<-1$ there appears to be a (at face value) smooth transition to a dynamically hot component."986" Similarly, in the upper panel of Figure 4 we see that the metal-rich component appears to have significant mean heliocentric velocity, while at [Fe/H] «-—1 one observes a smooth trend towards a relatively static velocity component."," Similarly, in the upper panel of Figure 4 we see that the metal-rich component appears to have significant mean heliocentric velocity, while at [Fe/H] $< -1$ one observes a smooth trend towards a relatively static velocity component."987" One caveat in this is that the metallicity errors are large, and we cannot rule our a bimodal distribution with the *wings' of the errors simulating a smoother transition between the two behaviors, although this would require a correlation between metallicity errors, measured radial velocity and velocity dispersion."," One caveat in this is that the metallicity errors are large, and we cannot rule our a bimodal distribution with the `wings' of the errors simulating a smoother transition between the two behaviors, although this would require a correlation between metallicity errors, measured radial velocity and velocity dispersion."988 The metal rich stars are best interpreted as part of the bar/bulge structure., The metal rich stars are best interpreted as part of the bar/bulge structure.989 Their kinematics show evidence of rotational support and bulk rotation and are consistent with data from the BRAVA survey in this region., Their kinematics show evidence of rotational support and bulk rotation and are consistent with data from the BRAVA survey in this region.990" The behavior of the more metal poor component, showing high velocity dispersion, and low to zero velocity relative to the Sun may be explained by a classical bulge or by inner halo stars."," The behavior of the more metal poor component, showing high velocity dispersion, and low to zero velocity relative to the Sun may be explained by a classical bulge or by inner halo stars."991" With a mean metal abundance of [Fe/H] ~—1.5 these stars appear to be best interpreted (at least provisionally) as an inner halo population, although bulges can of course be metal poor as well."," With a mean metal abundance of [Fe/H] $\sim -1.5$ these stars appear to be best interpreted (at least provisionally) as an inner halo population, although bulges can of course be metal poor as well."992" This is consistent with the observations by Zoccalietal.(2008) and Babusiauxetal.(2010), albeit for a single sightline."," This is consistent with the observations by \cite{zoccali08} and \cite{babusiaux10}, albeit for a single sightline."993" However, the BRAVA data show no classical bulge component fitting their dynamical model (Shen2010)."," However, the BRAVA data show no classical bulge component fitting their dynamical model \citep{shen10}."994". One possibility is that by selecting M giants and using the Calcium Triplet as their main radial velocity indicator, BRAVA may be biased against lower metallicity stars and therefore preferentially miss the high c component."," One possibility is that by selecting M giants and using the Calcium Triplet as their main radial velocity indicator, BRAVA may be biased against lower metallicity stars and therefore preferentially miss the high $\sigma$ component."995" The properties of galactic globular clusters present an interesting analogy with what is observed here: metal rich clusters, with mean [Fe/H] of ~—0.7 are believed to be associated with the bulge and are supported at least in part by rotation, whereas inner halo clusters have mean [Fe/H] of ~—1.6 their kinematics are dominated by random motions and at most very slow rotation."," The properties of galactic globular clusters present an interesting analogy with what is observed here: metal rich clusters, with mean [Fe/H] of $\sim -0.7$ are believed to be associated with the bulge and are supported at least in part by rotation, whereas inner halo clusters have mean [Fe/H] of $\sim -1.6$ their kinematics are dominated by random motions and at most very slow rotation."996" It is tempting to speculate that the two components we see in our data are analogous to the metal-poor and metal-rich globular clusters, whose properties they appear to share to some extent (cf.,"," It is tempting to speculate that the two components we see in our data are analogous to the metal-poor and metal-rich globular clusters, whose properties they appear to share to some extent (cf.,"997 Babusiauxet for a similar two-component model for the bulge)., \citealt{babusiaux10} for a similar two-component model for the bulge).998" etal.(1995);Zoccali(2003) and Clarksonetal. have argued, on the basis of isochrone fits to bulge globular clusters and field stars, that the bulge formed nearly coevally with the halo."," \cite{ortolani95,zoccali03} and \cite{clarkson08} have argued, on the basis of isochrone fits to bulge globular clusters and field stars, that the bulge formed nearly coevally with the halo."999 Most globular clusters in the inner halo formed within +1 Gyr of each other (Marin-Franchetal.2009)., Most globular clusters in the inner halo formed within $\pm 1$ Gyr of each other \citep{marin09}.1000". If this applies to bulge stars as well, it would imply a rapid star formation process, at least for the inner regions («20 kpc) of the Milky Way If this is the case, the smooth transition between the metal-rich and metal-poor subsystems, with a ‘turn-over’ point at [Fe/H] ~—1 may imply that the bulge and halo components are continuous and that there is no clear dichotomy between the two (modulo the large errors in metal abundance)."," If this applies to bulge stars as well, it would imply a rapid star formation process, at least for the inner regions $ < 20$ kpc) of the Milky Way If this is the case, the smooth transition between the metal-rich and metal-poor subsystems, with a `turn-over' point at [Fe/H] $\sim -1$ may imply that the bulge and halo components are continuous and that there is no clear dichotomy between the two (modulo the large errors in metal abundance)."1001 This would be consistent with the BRAVA result that the bulge was formed (in a dynamical sense) from secular evolution at high redshift., This would be consistent with the BRAVA result that the bulge was formed (in a dynamical sense) from secular evolution at high redshift.1002" As long as the stars also formed rapidly, the observed a-element enhancements are not in disagreement with this hypothesis."," As long as the stars also formed rapidly, the observed $\alpha$ -element enhancements are not in disagreement with this hypothesis."1003exposure (red/grey line in Fig. 1)),exposure (red/grey line in Fig. \ref{TW_Hya_lc_34}) )1004 and with separate third order polynomials for each step in the 380 nm exposure (red/grey line in Fig. 2)), and with separate third order polynomials for each step in the 380 nm exposure (red/grey line in Fig. \ref{TW_Hya_lc_38}) )1005 to remove trends before further analysis., to remove trends before further analysis.1006" We use the data from the central fiber only, which contains essentially all the signal."," We use the data from the central fiber only, which contains essentially all the signal."1007" The lightcurve shows some drop-outs, where no data was read out and a strong increase towards the end of the night."," The lightcurve shows some drop-outs, where no data was read out and a strong increase towards the end of the night."1008" We selected a clean, continuous region of about 40 min before the twilight sets in for analysis."," We selected a clean, continuous region of about 40 min before the twilight sets in for analysis."1009" This part of the lightcurve is shown in Fig. 3,,"," This part of the lightcurve is shown in Fig. \ref{AATau_lc},"1010 binned to 1 s. The second line shows the evolution of the full width at half maximum (FWHM) as fitted to a comparison star in the accompanying CCD images (timing resolution 10 s)., binned to 1 s. The second line shows the evolution of the full width at half maximum (FWHM) as fitted to a comparison star in the accompanying CCD images (timing resolution 10 s).1011 Times with large FWHM always correspond to low flux because light is lost from the instrument., Times with large FWHM always correspond to low flux because light is lost from the instrument.1012" We cannot correct for this effect by means of a comparison star, because the CCD requires an exposure time of 10 s, which is much longer than the sampling rate of the target."," We cannot correct for this effect by means of a comparison star, because the CCD requires an exposure time of 10 s, which is much longer than the sampling rate of the target."1013 The lightcurve is mean subtracted before further analysis., The lightcurve is mean subtracted before further analysis.1014 We determine the typical timescale of fluctuations by the autocorrelation function., We determine the typical timescale of fluctuations by the autocorrelation function.1015 To search for periodicity we first calculated Lomb-Scargle periodograms of our lightcurves., To search for periodicity we first calculated Lomb-Scargle periodograms of our lightcurves.1016" Additionally we performed a wavelet analysis, because periodic signals with changing frequency would be washed out in a simple periodogram."," Additionally we performed a wavelet analysis, because periodic signals with changing frequency would be washed out in a simple periodogram."1017 We quantify the significance of our non-detection by analytic considerations and Monte-Carlo simulations., We quantify the significance of our non-detection by analytic considerations and Monte-Carlo simulations.1018 The autocorrelation function is defined as where X; are the Xxvalues of the de-trended and mean-subtracted lightcurve and Ν is the total number of points., The autocorrelation function is defined as where $X_j$ are the values of the de-trended and mean-subtracted lightcurve and $N$ is the total number of points.1019 L is the scale of the correlation., $L$ is the scale of the correlation.1020" A correlation coefficient of 1 indicates a perfect correlation, 0 the absence of a linear correlation and -1 a perfect anti-correlation."," A correlation coefficient of 1 indicates a perfect correlation, 0 the absence of a linear correlation and -1 a perfect anti-correlation."1021 We show the auto-correlation coefficients for TW Hya in Fig., We show the auto-correlation coefficients for TW Hya in Fig.1022 4 for both filters., \ref{TW_Hya_autocorr} for both filters.1023" The autocorrelation of scale 0 is always 1; in TW Hya the coefficient differs significantly from the coefficient calculated for a comparison star for 2-4 s. The autocorrelation function for AA Tau is different from zero for about 10 s, however we find similar results using an intrinsically stable white dwarf from a different observation as comparison object."," The autocorrelation of scale 0 is always 1; in TW Hya the coefficient differs significantly from the coefficient calculated for a comparison star for 2-4 s. The autocorrelation function for AA Tau is different from zero for about 10 s, however we find similar results using an intrinsically stable white dwarf from a different observation as comparison object."1024 Apparently 10 s is the, Apparently 10 s is the1025bevond the light evlinder. contributing to peak 1 of the high-cnerey5 profile.,"beyond the light cylinder, contributing to peak 1 of the high-energy profile."1026 Hence. strong5 racio 1pulses with more pronounced. precursors. whieh correspond to large E. should be accompanied by stronger5 high-cnerey5S. emission. at. the position of peak 1 and weaker emission in the trough.," Hence, strong radio pulses with more pronounced precursors, which correspond to large $\Gamma$, should be accompanied by stronger high-energy emission at the position of peak 1 and weaker emission in the trough."1027 This also agrees with the observed. trend., This also agrees with the observed trend.1028 We have considered. the processes of. spontaneous. and induced. scattering olf the particles performing relativistic helical motion in an external magnetic field., We have considered the processes of spontaneous and induced scattering off the particles performing relativistic helical motion in an external magnetic field.1029 The theory is applied to the radio wave scattering olf the secondary plasma particles in the pulsar magnetosphere., The theory is applied to the radio wave scattering off the secondary plasma particles in the pulsar magnetosphere.1030 The particles are believed. to have substantial gvration energies due to resonant absorption of the radio emission in the outer magnetosphere., The particles are believed to have substantial gyration energies due to resonant absorption of the radio emission in the outer magnetosphere.1031 In application to the Vela pulsar. the induced scattering of radio waves between the states well below the resonance. wy=wu9. can be cllicient.," In application to the Vela pulsar, the induced scattering of radio waves between the states well below the resonance, $\omega\eta=\omega^\prime\eta^\prime\ll\Omega$, can be efficient."1032 An extremely bright anc narrow radio beam is scattered. chiellv into the background. in the direction corresponding to the maximum. scattering probability.," An extremely bright and narrow radio beam is scattered chiefly into the background, in the direction corresponding to the maximum scattering probability."1033 The scattered. component is directed approximately along the ambient magnetic feld. A.L/s |. and appears in the radio profile as a precursor to the main pulse.," The scattered component is directed approximately along the ambient magnetic field, $\theta_\mathrm{max}^\prime\sim 1/\gamma_\Vert$ , and appears in the radio profile as a precursor to the main pulse."1034 This scenario for the first time allows to explain the main features of the radio profile structure of the Vela pulsar ancl its pulse-to-pulse Iuctuations., This scenario for the first time allows to explain the main features of the radio profile structure of the Vela pulsar and its pulse-to-pulse fluctuations.1035 In the case considered. the induced scattering transfers the main pulse intensity to the higher frequencies. a’~wnt w.," In the case considered, the induced scattering transfers the main pulse intensity to the higher frequencies, $\omega^\prime\sim\omega\theta^2\gamma_\Vert^2\gg\omega$ ."1036 With the decreasing spectrum. of the pulsar radio emission. this implies intensity enhancement. of the radio profile at a fixed frequency.," With the decreasing spectrum of the pulsar radio emission, this implies intensity enhancement of the radio profile at a fixed frequency."1037 Thus. larger scattering ellieiencies should result in both higher pulse intensities an more pronounced precursors.," Thus, larger scattering efficiencies should result in both higher pulse intensities and more pronounced precursors."1038 Such a correlation is really observed in the Vela pulsar (Ixrishnamohan&Downs1983)., Such a correlation is really observed in the Vela pulsar \citep{kd83}.1039. Aloreover. stronger scattering means that a larger part. of the main pulse can be scattered: ellicientlv. the region of substantial suppression extending from the very trailing edge of the profile toward earlier phases.," Moreover, stronger scattering means that a larger part of the main pulse can be scattered efficiently, the region of substantial suppression extending from the very trailing edge of the profile toward earlier phases."1040 As the scattering of the leading edge of the main pulse makes the dominan contribution to the precursor formation. higher scattering elliciencies imply not only larger intensities but. also earlier phases of the precursor peak.," As the scattering of the leading edge of the main pulse makes the dominant contribution to the precursor formation, higher scattering efficiencies imply not only larger intensities but also earlier phases of the precursor peak."1041 The observations do reveal carly arrival of the precursor component in strong. pulses (Ixrishnuzunohan&Downs1983)., The observations do reveal early arrival of the precursor component in strong pulses \citep{kd83}.1042. The racio pulse is also subject to spontaneous scattering by the spiraling particles., The radio pulse is also subject to spontaneous scattering by the spiraling particles.1043 The photons below the resonance. we)9. are chiclly scattered to high. harmonics of the particle evrolrequency. s~55.," The photons below the resonance, $\omega\eta\ll\Omega$, are chiefly scattered to high harmonics of the particle gyrofrequency, $s\sim\gamma_0^3$."1044 We have analyzed the characteristics of the scattered: radiation. in detail and compared them with those of the svnchrotron. radiation of the same particle., We have analyzed the characteristics of the scattered radiation in detail and compared them with those of the synchrotron radiation of the same particle.1045 In. particular. it is found that the scattered power peaks at somewhat higher energies than the svnehrotron one.," In particular, it is found that the scattered power peaks at somewhat higher energies than the synchrotron one."1046 In application to the Vela pulsar. the estimates of the spectral maxima vield ~0.8 keV. anc ~0.Zac keV. respectively.," In application to the Vela pulsar, the estimates of the spectral maxima yield $\sim 0.8$ keV and $\sim 0.2$ keV, respectively."1047 The svnchrotron luminosity is 1 + and well agrees with the observed values in the soft. N-rav. band. (e.gPavlovetal.2001)., The synchrotron luminosity is $\sim 10^{31}$ $^{-1}$ and well agrees with the observed values in the soft X-ray band \citep[e.g][]{pavlov_01}.1048.. Although the total luminosity provided by the scattering is less. in the range bevond the svnchrotron maximum the scattered power may still contribute substantially.," Although the total luminosity provided by the scattering is less, in the range beyond the synchrotron maximum the scattered power may still contribute substantially."1049 Ht is important to note an extremely strong dependence of the scattered: power on the Lorentz-factor of the particle gvration. La;x55. which means à strong relation of the scattered. component to the racio intensity.," It is important to note an extremely strong dependence of the scattered power on the Lorentz-factor of the particle gyration, $L_\mathrm{sc}\propto\gamma_0^6$, which means a strong relation of the scattered component to the radio intensity."1050" Similarly to the svnchrotron emission. the scattered radiation concentrates close to the ambient magnetic field direction. 0,-1.po "," Similarly to the synchrotron emission, the scattered radiation concentrates close to the ambient magnetic field direction, $\theta^\prime\sim 1/\gamma_\Vert$."1051Given that (the. synchrotron. and scattered. emissions originate well above the radio emission region and inside the light evlinder. they should. precede the main radio pulse by =307.," Given that the synchrotron and scattered emissions originate well above the radio emission region and inside the light cylinder, they should precede the main radio pulse by $\la 30^\circ$."1052 This position can be identilied with that of peaks 3 and 4 of the soft N-ray. profile. which urn into the trough at somewhat higher energies (e.g.Larel-ingetal.2002).," This position can be identified with that of peaks 3 and 4 of the soft X-ray profile, which turn into the trough at somewhat higher energies \citep[e.g.][]{h02}."1053. I£ the svnchrotron re-emission. continues xvond the light exlinder. it may contribute to peak 1 of the »olile.," If the synchrotron re-emission continues beyond the light cylinder, it may contribute to peak 1 of the profile."1054 The radio emission of the Vela pulsar is believed. to xwticipate in both the spontaneous and induced scatterings. and the interplay between these processes can account or the observed. X-ray radio correlation.," The radio emission of the Vela pulsar is believed to participate in both the spontaneous and induced scatterings, and the interplay between these processes can account for the observed X-ray – radio correlation."1055 If the main oulse is less suppressed by the induced scattering to the oecursor. it is more cllicicntly scattered to high. energies.," If the main pulse is less suppressed by the induced scattering to the precursor, it is more efficiently scattered to high energies."1056 This is consistent with the observations: weaker racio pulses with less pronounced. precursors are accompanied. by. the ügh-energv pulses with stronger emission in the trough (Lommenοἱal.2007)., This is consistent with the observations: weaker radio pulses with less pronounced precursors are accompanied by the high-energy pulses with stronger emission in the trough \citep{l07}.1057. Given that the variations of the scattering elliciencies are determined by the [uctuations of sy. larger 5o implv less cllicient induced: scattering and simultaneously stronger high-energv. luminosities.," Given that the variations of the scattering efficiencies are determined by the fluctuations of $\gamma_0$, larger $\gamma_0$ imply less efficient induced scattering and simultaneously stronger high-energy luminosities."1058 For smaller συ the synchrotron re-emission. is expected to continue bevond the light evlinder. contributing to peak 1 of he high-enerew profile.," For smaller $\gamma_0$ the synchrotron re-emission is expected to continue beyond the light cylinder, contributing to peak 1 of the high-energy profile."1059 At the same time. smaller 55 imply more efficient. induced. scattering. à more pronounced. racio oecursor and stronger resultant radio pulses.," At the same time, smaller $\gamma_0$ imply more efficient induced scattering, a more pronounced radio precursor and stronger resultant radio pulses."1060 ALL this is in ine with the observed trends (Ixrishnamohan&Downs1983:Lommenetal. 2007).," All this is in line with the observed trends \citep{kd83,l07}."1061. Thus. our mocel explains the salient eatures of the radio profile formation of the Vela pulsar. the »eculiarities of its soft X-ray profile as well as the observed NM-ravo radio connection.," Thus, our model explains the salient features of the radio profile formation of the Vela pulsar, the peculiarities of its soft X-ray profile as well as the observed X-ray – radio connection."1062 Lt should be noted that our model of the radio profile ormation in the Vela pulsar is too simplified. since it includes only two components. the main pulse and precursor.," It should be noted that our model of the radio profile formation in the Vela pulsar is too simplified, since it includes only two components, the main pulse and precursor."1063 lrishnamohan&Downs(1983) have established. four components of the Vela's radio profile., \citet{kd83} have established four components of the Vela's radio profile.1064 Later on Johnstonetal.(2001). and Wrameretal.(2002). have. cliscovered sporadic activity at the leading edge of the pulse (the so-called. giant micro pulses) and in the bump region in the trailing part of the profile., Later on \citet{j01} and \citet{j02} have discovered sporadic activity at the leading edge of the pulse (the so-called giant micro pulses) and in the bump region in the trailing part of the profile.1065 Further development of our mocel is needed in order to include these phenomena., Further development of our model is needed in order to include these phenomena.1066 However. we believe that these. peculiarities do not alfect. the X-ray radio connection noticeably.," However, we believe that these peculiarities do not affect the X-ray – radio connection noticeably."1067 We have considered the induced. scattering in the approximation of a strong magnetic field., We have considered the induced scattering in the approximation of a strong magnetic field.1068 This process can x ellicient only i£ the incident ancl scattered waves have he ordinary polarization (the A-polarization). i.e. if their electric vectors are in the plane of the ambicnt magnetic ied.," This process can be efficient only if the incident and scattered waves have the ordinary polarization (the A-polarization), i.e. if their electric vectors are in the plane of the ambient magnetic field."1069 In the Vela pulsar. only one polarization mode is oesent.," In the Vela pulsar, only one polarization mode is present."1070 Various emission theories used to identify it with he ordinary modo., Various emission theories used to identify it with the ordinary mode.1071 The point is that the extraordinary. (D-) mode has the vacuum dispersion. and its direct. generation o» any plasma mechanism seems. problematic.," The point is that the extraordinary (B-) mode has the vacuum dispersion, and its direct generation by any plasma mechanism seems problematic."1072 The recent ugh-cnerey observations of the Velas pulsar wind nebula, The recent high-energy observations of the Vela's pulsar wind nebula1073values.,values.1074 Since the position of on the CCD was altered between the exposures. this procedure produced a clean correction image consisting purely of the fringe pattern.," Since the position of on the CCD was altered between the exposures, this procedure produced a clean correction image consisting purely of the fringe pattern."1075 After correction for the fringe pattern. the individual frames were registered and summed.," After correction for the fringe pattern, the individual frames were registered and summed."1076 The summed image shown in Figure | has a FWHM of 0.80 aresee and the sky background ts flat within over the entire field of view.," The summed image shown in Figure \ref{kentta}1077 has a FWHM of 0.80 arcsec and the sky background is flat within over the entire field of view."1078 The field was calibrated using stars 3. 5. and 6 for which published Cousins I-band magnitudes.," The field was calibrated using stars 3, 5, and 6 for which published Cousins I-band magnitudes."1079 Unfortunately. all three stars were saturated in their cores (7 < 0744 - 0788). and could not be used directly for calibration purposes.," Unfortunately, all three stars were saturated in their cores $r$ $<$ 4 - 8), and could not be used directly for calibration purposes."1080 Instead. we used a 45 s 1-band exposure of the field acquired immediately before the sequence.," Instead, we used a 45 s i-band exposure of the field acquired immediately before the sequence."1081 The FWHM in this image is significantly worse than in the the sequence and stars 3. 5. and 6 are not saturated.," The FWHM in this image is significantly worse than in the the sequence and stars 3, 5, and 6 are not saturated."1082 Using these stars we first determined the zero point of the 45 s exposure., Using these stars we first determined the zero point of the 45 s exposure.1083 All three stars gave consistent results. the maximum difference between the zero points was 0.047 mag. consistent with the errors of the Ic magnitudes 0.04-0.05 nag in?.," All three stars gave consistent results, the maximum difference between the zero points was 0.047 mag, consistent with the errors of the $_{\rm C}$ magnitudes 0.04-0.05 mag in."1084. Using this zero point. we determined the Ic magnitudes of stars S1. S2. and S3 and from these three stars the zero point of the combined long exposure.," Using this zero point, we determined the $_{\rm C}$ magnitudes of stars S1, S2, and S3 and from these three stars the zero point of the combined long exposure."1085 The uncertainly in the zero point was dominated by the uncertainty in the Ic magnitudes of stars 3. 5. and 6 and by the color effects between our i-band filter and the Ic filter.," The uncertainly in the zero point was dominated by the uncertainty in the $_{\rm C}$ magnitudes of stars 3, 5, and 6 and by the color effects between our i-band filter and the $_{\rm C}$ filter."1086 We estimate the former contribution to the uncertainty to be 0.03 mag., We estimate the former contribution to the uncertainty to be 0.03 mag.1087 To estimate the magnitude of the color effects. we used the established relation between Ic and and SDSS 1 given by Lupton (2005) in the SDSS web pages: 72/-0.37806—0.3974.," To estimate the magnitude of the color effects, we used the established relation between $_{\rm C}$ and and SDSS i given by Lupton (2005) in the SDSS web pages: $I = i - 0.3780(i-z) -0.3974$ ."1088 The SDSS 1 has approximately the same width as our i-band filter. but the central wavelength is ~ 50 nm lower.," The SDSS i has approximately the same width as our i-band filter, but the central wavelength is $\sim$ 50 nm lower."1089 For stars 3. 5. and 6. the /—z color is  -0.01 — 0.06.," For stars 3, 5, and 6, the $i-z$ color is $\sim$ -0.01 – 0.06."1090 Assuming stars S1-S3 have similar colors and since an elliptical galaxy at z = 0.3-0.5 has a color /—z~0.4—0.5. the color effects between the calibration stars and the host galaxy of may cause an error of up to 0.2 mag in SDSS 1. Since our filter is close to SDSS 1. we expect the color effects to be similar in value.," Assuming stars S1-S3 have similar colors and since an elliptical galaxy at z = 0.3-0.5 has a color $i-z \sim10910.4-0.5$, the color effects between the calibration stars and the host galaxy of may cause an error of up to 0.2 mag in SDSS i. Since our filter is close to SDSS i, we expect the color effects to be similar in value."1092 Since the color errors dominate the total error of the zero point. we assign a formal error of 0.2 mag to the magnitude zero point.," Since the color errors dominate the total error of the zero point, we assign a formal error of 0.2 mag to the magnitude zero point."1093 The host galaxy was analyzed by fitting a two-dimensional surface brightness model to the light distribution of0716-714., The host galaxy was analyzed by fitting a two-dimensional surface brightness model to the light distribution of.1094.. Details of this process can be found in?., Details of this process can be found in.1095. In short. the model consists of two components. an unresolved core representing the BL Lac nucleus and a de Vaucouleurs profile representing the host galaxy.," In short, the model consists of two components, an unresolved core representing the BL Lac nucleus and a de Vaucouleurs profile representing the host galaxy."1096 Previous imaging has shown that the de Vaucouleurs profile describes well the surface brightness profiles of BL Lac host galaxies. although deviations from this law have also been observed?).," Previous imaging has shown that the de Vaucouleurs profile describes well the surface brightness profiles of BL Lac host galaxies, although deviations from this law have also been observed."1097. For our purposes. it is sufficient to know that BL Lac host galaxies are bulge-dominated systems and no BL Lacertae object has ever been reliably associated with a disk type host.," For our purposes, it is sufficient to know that BL Lac host galaxies are bulge-dominated systems and no BL Lacertae object has ever been reliably associated with a disk type host."1098" The model is described by 5 parameters: position (vv). core magnitude mic. host galaxy magnitude 77,5. and host galaxy effective radius ray."," The model is described by 5 parameters: position $x$ $y$ ), core magnitude $m_C$, host galaxy magnitude $m_H$, and host galaxy effective radius $r_{\rm eff}$."1099 The model parameters are adjusted using an iterative Levenberg-Marquardt. loop. until the minimum value of the y* statistic between the model and the data is found.," The model parameters are adjusted using an iterative Levenberg-Marquardt loop, until the minimum value of the $\chi^2$ statistic between the model and the data is found."1100 Only pixels within 955 from the center of were included in the fit and pixels influenced by an I = 20.8 mag nearby companion 4722 west of were excluded., Only pixels within 5 from the center of were included in the fit and pixels influenced by an I = 20.8 mag nearby companion 2 west of were excluded.1101 The model is convolved by the observed PSF. determined from stars in the vicinity of0716+714.," The model is convolved by the observed PSF, determined from stars in the vicinity of."1102. Since stars 5 and 6 were saturated in their cores (r « 0.4 arcsec: the peak counts of remained under 58 000 ADUS in all images.).," Since stars 5 and 6 were saturated in their cores (r $<$ 0.4 arcsec; the peak counts of remained under 58 000 ADUs in all images.),"1103 we constructed the PSF from stars 5 and SI by joining smoothly the outer parts of star 5 with the inner part of SI (hereafter we refer to this PSF as PSFI)., we constructed the PSF from stars 5 and S1 by joining smoothly the outer parts of star 5 with the inner part of S1 (hereafter we refer to this PSF as PSF1).1104 A second PSF was constructed in à similar way from stars 6 and S2 (PSF2) to study the effect of PSF variability over the field of view., A second PSF was constructed in a similar way from stars 6 and S2 (PSF2) to study the effect of PSF variability over the field of view.1105 Since the host galaxy is faint compared to the nucleus. the results are highly dependent on the accuracy of the PSF and its stability over the field of view.," Since the host galaxy is faint compared to the nucleus, the results are highly dependent on the accuracy of the PSF and its stability over the field of view."1106 We analyzed the PSF stability by extracting the surface brightness profiles of stars 2. 3. 5. 6. SI. and S2 and scaling them to the same magnitude.," We analyzed the PSF stability by extracting the surface brightness profiles of stars 2, 3, 5, 6, S1, and S2 and scaling them to the same magnitude."1107 In the lower panel of Fig. 2..," In the lower panel of Fig. \ref{prof},"1108 we show the surface brightness profiles of these stars relative to star 5. which is equal to PSFI at 7> 0744.," we show the surface brightness profiles of these stars relative to star 5, which is equal to PSF1 at $r >$ 4."1109 We note that in this representation star 5 appears as a horizontal line., We note that in this representation star 5 appears as a horizontal line.1110 The surface brightness profiles of the stars agree to within 0.2 mag all the way to the outer fitting radius of 9755., The surface brightness profiles of the stars agree to within 0.2 mag all the way to the outer fitting radius of 5.1111 The rms scatter between the surface brightness profiles is close to the center and increases to ~ at the outer fitting radius., The rms scatter between the surface brightness profiles is close to the center and increases to $\sim$ at the outer fitting radius.1112Lt is well established that the measurement of the galaxy Iuminosity function (LE) is sensitive to the tvpe of galaxy being sampled.,It is well established that the measurement of the galaxy luminosity function (LF) is sensitive to the type of galaxy being sampled.1113 \lorphologically earlv-type galaxies tend to be systematically brighter than their late-type counterparts. resulting in LE estimates with signilicantlv brighter AZ* arc shallower faint end slope a (e.g. Efstathiou. Ellis Peterson LOSS: Loveday et al.," Morphologically early-type galaxies tend to be systematically brighter than their late-type counterparts, resulting in LF estimates with significantly brighter $M^*$ and shallower faint end slope $\alpha$ (e.g. Efstathiou, Ellis Peterson 1988; Loveday et al."1114 1992: Blanton et al., 1992; Blanton et al.1115 2001)., 2001).1116 Similar trends are also present if one selects galaxies basec on Ho equivalent widths (Lovedav. Tresse Macdox. 1999). OL equivalent widths (Ellis et al.," Similar trends are also present if one selects galaxies based on $\alpha$ equivalent widths (Loveday, Tresse Maddox 1999), [OII] equivalent widths (Ellis et al."1117 1996) or colour (Lin e al., 1996) or colour (Lin et al.1118 1906: Alarzke Da Costa 1997)., 1996; Marzke Da Costa 1997).1119 Understanding. anc being able to quantify this variation in the LE is of grea importance to a Full understanding of galaxy formation anc evolution., Understanding and being able to quantify this variation in the LF is of great importance to a full understanding of galaxy formation and evolution.1120 Up until now the greatest obstacle. faced. when attempting to determine how the LE varied. with dilleren galaxy properties was the fact that one was required. to divide already small cata sets. thereby. losing much of the statistical significance in the LE estimations.," Up until now the greatest obstacle faced when attempting to determine how the LF varied with different galaxy properties was the fact that one was required to divide already small data sets, thereby losing much of the statistical significance in the LF estimations."1121 In this paper we address this issue by making use of a subset of the ealaxies observed to date in the 2db Galaxy Redshift Survey (2PdEORS Maddox: 1998: Colless 2001)., In this paper we address this issue by making use of a subset of the galaxies observed to date in the 2dF Galaxy Redshift Survey (2dFGRS Maddox 1998; Colless 2001).1122 This subset is the largest single data set used in the calculation of LEs and can easily be divided several times whilst still maintaining very precise statistics., This subset is the largest single data set used in the calculation of LFs and can easily be divided several times whilst still maintaining very precise statistics.1123" The 24ECGHRS is a joint. Ulx-Xustralian elfort to map the clistribution of galaxies down to an extinction corrected b, maenituce of 19.45 (median redshift z2 0.1).", The 2dFGRS is a joint UK-Australian effort to map the distribution of galaxies down to an extinction corrected $\bj$ magnitude of $19.45$ (median redshift $z\simeq0.1$ ).1124 In so doing we expect to obtain 250.000 galaxy spectra from which redshifts can be determined.," In so doing we expect to obtain 250,000 galaxy spectra from which redshifts can be determined."1125 This is à factor of 10 more than any previous redshift survey., This is a factor of 10 more than any previous redshift survey.1126 1n order to divide our data set in a meaningful way we develop in this paper a classification of the galaxies based upon their observed spectra., In order to divide our data set in a meaningful way we develop in this paper a classification of the galaxies based upon their observed spectra.1127 As with other astronomical ata. there are different approaches for analysing galaxy spectra (see e.g. Lahay 2000 and references therein)," As with other astronomical data, there are different approaches for analysing galaxy spectra (see e.g. Lahav 2000 and references therein)."1128 If one has a well defined physical model for galaxy. spectra. then it is appropriate to estimate parameters of interest. (such as age and star-formation rate) e.g. by Maximum Likelihood. irectly using all the spectral. bins. or via à compressec version of the data designed to give maximum information on the physical parameters of interest (c.g. Leavens. Jimenez Lahay 2000).," If one has a well defined physical model for galaxy spectra, then it is appropriate to estimate parameters of interest (such as age and star-formation rate) e.g. by Maximum Likelihood, directly using all the spectral bins, or via a compressed version of the data designed to give maximum information on the physical parameters of interest (e.g. Heavens, Jimenez Lahav 2000)."1129 1. on the other hand. one prefers to Ie the data ‘speak for themselves. in a. model-independen wav. then it is more useful to look at the distribution of the galaxies in the high-dimensional space defined. by. the spectral bins.," If, on the other hand, one prefers to let the data `speak for themselves', in a model-independent way, then it is more useful to look at the distribution of the galaxies in the high-dimensional space defined by the spectral bins."1130 Lt is then possible to either look for distinc eroups. e.g. early and late types. or more refined classes (e.g. Slonim et al.," It is then possible to either look for distinct groups, e.g. early and late types, or more refined classes (e.g. Slonim et al."1131 2000)., 2000).1132 Xn alternative. which we present here. is to find a continuous (sequence-ike) parameterisation of the spectral features. which can later be divided into subsets.," An alternative, which we present here, is to find a continuous (sequence-like) parameterisation of the spectral features, which can later be divided into subsets."1133 The parameterisation. we develop. is. denoted. by a and essentially represents a measure of the average absorption or emission line strength present in cach ealaxy’s spectrum., The parameterisation we develop is denoted by $\eta$ and essentially represents a measure of the average absorption or emission line strength present in each galaxy's spectrum.1134 This classification is robust to the known instrumental uncertainties and also has the advantage of being casily interpreted. in terms of the current star-forniation. present in cach galaxy., This classification is robust to the known instrumental uncertainties and also has the advantage of being easily interpreted in terms of the current star-formation present in each galaxy.1135 The set of galaxies we use in this paper comprises a sample of 75.000 galaxy spectra.," The set of galaxies we use in this paper comprises a sample of 75,000 galaxy spectra."1136 This is more than 12 times larger than the one used in our previous analysis (Folkes e al., This is more than 12 times larger than the one used in our previous analysis (Folkes et al.1137 1999. hereafter E99) and is by far the largest sample usec to date.," 1999, hereafter F99) and is by far the largest sample used to date."1138 As well as showing the latest LE determinations we wil also present à new procedure for caleulating self-consisten Áx-corrections and investigate possible fibre-aperture biases., As well as showing the latest LF determinations we will also present a new procedure for calculating self-consistent $K$ -corrections and investigate possible fibre-aperture biases.1139 The outline of this paper is as follows., The outline of this paper is as follows.1140" Section 2 briefly summarises the 2dFCRS and describes the cata se we are using. Section 3. outlines our method for dividing the data set based: on our parameterisation of the galaxy spectral type. η,"," Section \ref{section:data} briefly summarises the 2dFGRS and describes the data set we are using, Section \ref{section:spectypes} outlines our method for dividing the data set based on our parameterisation of the galaxy spectral type, $\eta$."1141 Section 4. gives a detailed description of the calculations involved in estimating the LE and in Section we discuss our results and future work., Section \ref{section:lumfunctions} gives a detailed description of the calculations involved in estimating the LF and in Section \ref{section:discussion} we discuss our results and future work.1142 The 2dEGIS has already observed: approximately: 200.000 unique galaxies for which it has obtained redshifts.," The 2dFGRS has already observed approximately 200,000 unique galaxies for which it has obtained redshifts."1143 The survey. once complete. will cover approximately 2000717 on the sky. split between two independent strips: one in the northern Galactic hemisphere and the other centred roughly on the south Galactic pole.," The survey, once complete, will cover approximately $\sq\degr$ on the sky, split between two independent strips: one in the northern Galactic hemisphere and the other centred roughly on the south Galactic pole."1144 In addition to this there are 99 random fields in the southern Galactic cap., In addition to this there are 99 random fields in the southern Galactic cap.1145 The 2dE instrument itself is capable of observing up to 400 ealaxy spectra simultaneously (Lewis et al..," The 2dF instrument itself is capable of observing up to 400 galaxy spectra simultaneously (Lewis et al.,"1146. 2001)., 2001).1147 Each galaxy has been selected [rom a revised and extended version of the Automatecl Plate Measuring. (APA) galaxy catalogue (Macdox et al., Each galaxy has been selected from a revised and extended version of the Automated Plate Measuring (APM) galaxy catalogue (Maddox et al.1148 1990) in order to determine its position and magnitude., 1990) in order to determine its position and magnitude.1149 It is then observed as part of the survey through a 40s; (2 arcsec) diameter. optical libre., It is then observed as part of the survey through a $\mu$ m $\sim2$ arcsec) diameter optical fibre.1150 The 5b; magnitudes we use in this analysis are total magnitudes derived. from. updated: versions of the original APM scans., The $b_{\rm J}$ magnitudes we use in this analysis are total magnitudes derived from updated versions of the original APM scans.1151 These magnitudeso have been updated to take into account new CCD calibration data (see Colless ot al., These magnitudes have been updated to take into account new CCD calibration data (see Colless et al.1152 2001 and Norbere et al..," 2001 and Norberg et al.,"1153 2001) and are believed to have an ros error of approximately 0.15 magnitudes., 2001) and are believed to have an rms error of approximately $\pm0.15$ magnitudes.1154 Each of the spectra observed. using the 2d instrument spans 1024 channels with a spectral scale of 4.3A per pixel: the EPNCLEM is measured by are lines to be of the order of 1.8- pixels., Each of the spectra observed using the 2dF instrument spans 1024 channels with a spectral scale of $\rm{\AA}$ per pixel; the FWHM is measured by arc lines to be of the order of 1.8-2.5 pixels.1155 Typically at the survey limit the observed spectra have an average signal-to-noise ratio of 10 per pixel. sullicient for determining redshifts ancl performing spectral analyses.," Typically at the survey limit the observed spectra have an average signal-to-noise ratio of $\sim10$ per pixel, sufficient for determining redshifts and performing spectral analyses."1156 For the purposes of our spectral analysis we have restricted: ourselves to the redshift range of 0.01<z«0.2 so that we are left with a uniform sample of galaxy spectra in the rest-frame wavelength range spanning OL] to Ho., For the purposes of our spectral analysis we have restricted ourselves to the redshift range of $0.01<z<0.2$ so that we are left with a uniform sample of galaxy spectra in the rest-frame wavelength range spanning [OII] to $\alpha$.1157 A detailed outline of the spectral reduction. pipeline is given in Colless et al. (, A detailed outline of the spectral reduction pipeline is given in Colless et al. (11582001) and E99 which we follow with only minor moclifications.,2001) and F99 which we follow with only minor modifications.1159 The data that we have available for this analysis include all observations up until January 2001., The data that we have available for this analysis include all observations up until January 2001.1160 At this stage à total, At this stage a total1161The CHIPS data show that the EUV iron lines are extremely faint. and even the fIuxes analvzed here are likely (ο include a local / foreground component. leaving even less room for interstellar flux.,"The CHIPS data show that the EUV iron lines are extremely faint, and even the fluxes analyzed here are likely to include a local / foreground component, leaving even less room for interstellar flux."1162 Although previous works have suggested that the Fe features are weaker than predicted by collisional ionization equilibrium mocdels. some of those results have rested on uncertain and untestable assumptions.," Although previous works have suggested that the Fe features are weaker than predicted by collisional ionization equilibrium models, some of those results have rested on uncertain and untestable assumptions."1163 Using a point-source spectrometer aboard the Extreme Ultraviolet Explorer. Jelinskvetal.(1995) set an upper limit to the emission measure near 1077 IX that is also lower than expected for the local hot bubble (but not as light as our limits).," Using a point-source spectrometer aboard the Extreme Ultraviolet Explorer, \citet{JVE95} set an upper limit to the emission measure near $^{5.8}$ K that is also lower than expected for the local hot bubble (but not as tight as our limits)."1164" However. owing in part to the low resolution of the instrument. it was difficult to confirm that the very large background in the EUVE signal could be subtracted precisely,"," However, owing in part to the low resolution of the instrument, it was difficult to confirm that the very large background in the EUVE signal could be subtracted precisely."1165 Blochetal.(2002) interpret ALENIS multilaver imaging of the difDuse background wilh a model in which an astrophysical EUY {hax is assumed (o scale with (he Wisconsin map. and conclude that (he Fe emission is no more Chan about 1/10 of the predicted Πας level.," \citet{Bloch02} interpret ALEXIS multilayer imaging of the diffuse background with a model in which an astrophysical EUV flux is assumed to scale with the Wisconsin B-band map, and conclude that the Fe emission is no more than about 1/10 of the predicted flux level."1166 The pulse-height distribution lor the Be-bancl proportional counter Wisconsin rocket data has been reported to require significant. depletions (i.e.. the pulse height distribution is inconsistent with full-strength iron lines) (Bloch1938).," The pulse-height distribution for the Be-band proportional counter Wisconsin rocket data has been reported to require significant depletions (i.e., the pulse height distribution is inconsistent with full-strength iron lines) \citep{B88}."1167. Using a rocket-borne calorimeter. MeCanunonetal.(2002) report a detection of emission from FeEX.X. and (unresolved) with a combined flux of about 100 LU. significantly greater than allowed by the CIIIPS data.," Using a rocket-borne calorimeter, \citet{MCC02} report a detection of emission from Fe, and (unresolved) with a combined flux of about 100 LU, significantly greater than allowed by the CHIPS data."1168 The calorimeter signal (4 events) exceeds (he expected background (0.3 events) and sets a lower limit of about 40 LU on the combined line strengths. a limit well above the best-fit CHIPS results.," The calorimeter signal (4 events) exceeds the expected background (0.3 events) and sets a lower limit of about 40 LU on the combined line strengths, a limit well above the best-fit CHIPS results."1169 Whether the dillerence in results arises from instrument. calibration effects. a statistical oddity. or genuine dillerences in fields observed (due to astrophysical or οσα] effects) is not vet clear.," Whether the difference in results arises from instrument calibration effects, a statistical oddity, or genuine differences in fields observed (due to astrophysical or local effects) is not yet clear."1170 In any case. the CIIIPS observations set robust constraints on important individual spectral lines.," In any case, the CHIPS observations set robust constraints on important individual spectral lines."1171 The CHIPS results are al variance with X-ray measurements interpreted. wilh similar nodels and subject to the same assumptions (solar abundances. collisional ionization equilibrium).," The CHIPS results are at variance with X-ray measurements interpreted with similar models and subject to the same assumptions (solar abundances, collisional ionization equilibrium)."1172 As noted above. temperatures greater than LO’? IX. and emission measures an order of nagnitude higher than the upper limits of Figure 3. have been reported.," As noted above, temperatures greater than $^{6.0}$ K, and emission measures an order of magnitude higher than the upper limits of Figure 3, have been reported."1173 Depletion of iron relative to other heavy elements generally helps reconcile the EUV and X-ray results. but the depletions reported previously have been only moderate: roughly a [actor of 3 reduction in (he gas phase abundance of elements including iron (Sandersetal.2001;al. 2002).," Depletion of iron relative to other heavy elements generally helps reconcile the EUV and X-ray results, but the depletions reported previously have been only moderate: roughly a factor of 3 reduction in the gas phase abundance of elements including iron \citep{Setal01,MCC02}."1174. The CHIPS results are qualitatively consistent will a poster presented by (2003).. concluding that a combination of depletion and a lower temperature (107? I) in (he local hot gas provides a better fit to the broadband X-ray. colors compared (to wnclepletect. roller models.," The CHIPS results are qualitatively consistent with a poster presented by \citet{Bellm03}, concluding that a combination of depletion and a lower temperature $^{5.85}$ K) in the local hot gas provides a better fit to the broadband X-ray colors compared to undepleted, hotter models."1175 Foreground absorption alone cannot reconcile (he CHIPS and X-ray results., Foreground absorption alone cannot reconcile the CHIPS and X-ray results.1176 A foreground, A foreground1177fit (he Nice point distribution.,fit the $N_{\rm IGC}$ point distribution.1178 The best fit to a Ixing distribution produces r;/r.=81.5 and ace—9.5x107., The best fit to a King distribution produces $r_t/r_c=81.5$ and $\hat\epsilon_{\rm IGC}=9.5\times 10^{-5}$.1179 The result has been plotted in the lower panel of Fig. 5.., The result has been plotted in the lower panel of Fig. \ref{massdis2}.1180 The reduced chi-square for the best fit results \2=0.84 which indicates that the fit is good enough., The reduced chi-square for the best fit results $\chi^2_{\nu}=0.84$ which indicates that the fit is good enough.1181 On the other hand. \7=3.0 for the r;/r.=14.2 king model fit.," On the other hand, $\chi^2_{\nu}=3.0$ for the $r_t/r_c=14.2$ king model fit."1182" For v=19 (as in our case). the integral probability of the \2 distribution is 0.001 for AZ>2.31. showing that (her,/r,.=14.2 king model fit can be rejected with a high confidence level."," For $\nu=19$ (as in our case), the integral probability of the $\chi^2_{\nu}$ distribution is 0.001 for $\chi^2_{\nu} \geq 2.31$, showing that the $r_t/r_c=14.2$ king model fit can be rejected with a high confidence level."1183" The concentration parameter r,/r.=81.5. implies a core radius of r.=154."," The concentration parameter $r_t/r_c=81.5$, implies a core radius of $r_c=154\arcsec$."1184 This is ol the order οἱ. or smaller than. the size of the Coma central galaxy NGC 4374.," This is of the order of, or smaller than, the size of the Coma central galaxy NGC 4874."1185" In [act. in Mavin-Franeh&Aparicio(2002).. the spatial structure of the NGC 4874 GCS was studied out (o a distance of 161.4"" from the galaxy. center. and its edge was probably not reached."," In fact, in \citet{MA02}, the spatial structure of the NGC 4874 GCS was studied out to a distance of $161.4\arcsec$ from the galaxy center, and its edge was probably not reached."1186 This suggests that the obtained Nqc;c excess around the center of Coma is produced by the outermost extension of the NGC 4874 primitive GCS rather than being a trace of an IGC svstem., This suggests that the obtained $N_{\rm IGC}$ excess around the center of Coma is produced by the outermost extension of the NGC 4874 primitive GCS rather than being a trace of an IGC system.1187 In other words. (he spatial distribution of Nye shown in Fig.," In other words, the spatial distribution of $N_{\rm IGC}$ shown in Fig."1188 5 is probably the map of the NGC 4874 GCS., \ref{massdis2} is probably the map of the NGC 4874 GCS.1189 This result suggests that an [GC population does not exist in Coma., This result suggests that an IGC population does not exist in Coma.1190 This points towarel GCs having been formed only. or almost only. [rom protogalactic clouds.," This points toward GCs having been formed only, or almost only, from protogalactic clouds."1191 None of them. or perhaps very few. could have formed in isolated regions.," None of them, or perhaps very few, could have formed in isolated regions."1192 This is in good agreement with the formation scenario for the GCSs (Blakesleeetal.1997:Forbes.Drodie.&Grillmairllarrisοἱal.1998:MeLaughlin. 1999).," This is in good agreement with the formation scenario for the GCSs \citep{B97,F97,H98,M99}."1193. If the distribution of IGCs followed the distribution of dark matter in galaxy clusters (Muzzio1987;West1993).. it. would be reasonable to expect that Nige would show a spatially extended. gradient. across de galaxy. cluster.," If the distribution of IGCs followed the distribution of dark matter in galaxy clusters \citep{M87,W93}, it would be reasonable to expect that $N_{\rm IGC}$ would show a spatially extended gradient across de galaxy cluster."1194 This not being (he case. il seems inappropriate to support the relationship between IGC and dark malter distributions. although it is true that the relationship could still exist but not be strong enough to be detected.," This not being the case, it seems inappropriate to support the relationship between IGC and dark matter distributions, although it is true that the relationship could still exist but not be strong enough to be detected."1195 On the other hand. a number of authors have argued. that accretion of IGCs might influence the formation and evolution of galaxies (Muzzio1987:White1937:al.1995:Cotéet 2001).," On the other hand, a number of authors have argued that accretion of IGCs might influence the formation and evolution of galaxies \citep{M87,W87,W95,C01}."1196. Our results indicate that accretion of IGCs is nota significant effect in galaxy formation and evolutionary processes in the Coma galaxies. since our conclusion is that IGCs do not exist in Coma.," Our results indicate that accretion of IGCs is nota significant effect in galaxy formation and evolutionary processes in the Coma galaxies, since our conclusion is that IGCs do not exist in Coma."1197" Gurzadvan&Mazure(2001) found the existence of three subgroups of galaxies in Coma. one of them associated) with the «D galaxy NGC 4874 and the other two with NGC 4889 and NGC 4839,"," \citet{GM01} found the existence of three subgroups of galaxies in Coma, one of them associated with the cD galaxy NGC 4874 and the other two with NGC 4889 and NGC 4839."1198 They conclude that the non-stationaritv of the dvnamical processes al work in the Coma core is due to the merging of small-scale groups of galaxies., They conclude that the non-stationarity of the dynamical processes at work in the Coma core is due to the merging of small-scale groups of galaxies.1199 Bi this context. each subgroup formed separately and then the merger between the different groups took place.," In this context, each subgroup formed separately and then the merger between the different groups took place."1200 In Marin-Franch&Aparicio (2002).. [ive galaxies belonging to the," In \citet{MA02}, , five galaxies belonging to the"

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