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Since hotspots have spectra with injection slope a ranging between 0.5 and 1 (as a reference. the classical value from the dilfuse particle acceleration at strong shocks is a=0.5. e.g. Meisenheimer et al. | Since hotspots have spectra with injection slope $\alpha$ ranging between 0.5 and 1 (as a reference, the classical value from the diffuse particle acceleration at strong shocks is $\alpha = 0.5$, e.g. Meisenheimer et al. |
1997). we decided. to consider the injection spectral index as a free parameter. | 1997), we decided to consider the injection spectral index as a free parameter. |
Such constraints allow us to determine the spectrum of the emitting electrons. (normalization. break ancl cut-oll energy). once the magnetic field strength has been assumed. and to calculate. the emission. from. either. svnchrotron-sel-Compton (SSC) or inverse-Compton scattering of the cosmic background radiation (1C-CM) expected from the hotspot (or jet) region (followingBrunettietal.2002). | Such constraints allow us to determine the spectrum of the emitting electrons (normalization, break and cut-off energy), once the magnetic field strength has been assumed, and to calculate the emission from either synchrotron-self-Compton (SSC) or inverse-Compton scattering of the cosmic background radiation (IC-CMB) expected from the hotspot (or jet) region \citep[following][]{gb02}. |
. Models described in. Brunettietal.(2002) ake also into account the boosting effects arising from a hotspot/jet that is moving at relativistic speeds and. oriented. at à. given angle with respect to our line of In Figures 4. to 6 we show the SED from the racio band to high energv emission. measured. for the hotspot components of 1105 South. together with the mocel fits. | Models described in \citet{gb02} take also into account the boosting effects arising from a hotspot/jet that is moving at relativistic speeds and oriented at a given angle with respect to our line of In Figures \ref{fig_spectra_105n} to \ref{fig_spectra_105s} we show the SED from the radio band to high energy emission measured for the hotspot components of 105 South, together with the model fits. |
Synchrotron models with an injection spectral index a=0.8 provide an adequate representation of the SLED of the central ancl southern components of 1105. South. with break frequencies ranging from 1077 to 1077 W/llz. while the eutolf frequencies are between 3 1043 and 107 W/Llz. | Synchrotron models with an injection spectral index $\alpha$ =0.8 provide an adequate representation of the SED of the central and southern components of 105 South, with break frequencies ranging from $\times$ $^{12}$ to $\times$ $^{13}$ W/Hz, while the cutoff frequencies are between $\times$ $^{14}$ and $\times$ $^{15}$ W/Hz. |
In both components. the upper limit to the X-ray emission does not allow us to constrain the validity of the SSC model (Figs. | In both components, the upper limit to the X-ray emission does not allow us to constrain the validity of the SSC model (Figs. |
5. and 6)). | \ref{fig_spectra_105c} and \ref{fig_spectra_105s}) ). |
On the other hand. the northern component of 1105. shows à. prominent X-ray emission. | On the other hand, the northern component of 105 shows a prominent X-ray emission. |
A svachrotron model (dashed. line in Fig. 4)) | A synchrotron model (dashed line in Fig. \ref{fig_spectra_105n}) ) |
may [it quite reasonably the racio. NIB. ancl X-ray emission. but it completely fails in reproducing the optical data. | may fit quite reasonably the radio, NIR and X-ray emission, but it completely fails in reproducing the optical data. |
An acleditional contribution of the SSC is not a viable option since it requires a magnetic field much smaller than hat obtained. assuming equipartition (see Section 5.4) (solid. lines). and implving an unreasonably large energy μισοί. | An additional contribution of the SSC is not a viable option since it requires a magnetic field much smaller than that obtained assuming equipartition (see Section 5.4) (solid lines), and implying an unreasonably large energy budget. |
On the other hand. the high energy emission is well modelled. by LO-CAIB (ο.Tavecchioetal.2000:Colottietal.POOL) where the CAIB photons are scattered w relativistic electrons. with Lorentz factor D6. and 6—5" with a magnetic field of 16 µ This model implies hat boosting effects play an important role in the X-ray emission. of this component. sugecsting that SI is more ikely à relativistic knot in the jet. rather than a hotspot eature. | On the other hand, the high energy emission is well modelled by IC-CMB \citep[e.g.][]{tavecchio00,celotti01} where the CMB photons are scattered by relativistic electrons with Lorentz factor $\Gamma \sim 6$, and $\theta$ $^{\circ}$ with a magnetic field of 16 $\mu$ G. This model implies that boosting effects play an important role in the X-ray emission of this component, suggesting that S1 is more likely a relativistic knot in the jet, rather than a hotspot feature. |
Phe weakness of this interpretation is that 1105 is à NLRG and its jets are expected to form a Large angle with our line of sight. | The weakness of this interpretation is that 105 is a NLRG and its jets are expected to form a large angle with our line of sight. |
Alternatively. the X-ray emission may ο svnchrotron from a different population of electrons. as sugeested in the case of the jet in 2273 (Jester et al. | Alternatively, the X-ray emission may be synchrotron from a different population of electrons, as suggested in the case of the jet in 273 (Jester et al. |
The analysis of the southern hotspot of 4445 as a single unresolved Component was carried. out. in. previous work bv Prietoctal.(2002):Macket(2009):Perlman (2010). | The analysis of the southern hotspot of 445 as a single unresolved component was carried out in previous work by \citet{aprieto02,mack09,perlman10}. |
. In this new analysis. the high spatial resolution and multiwavelength VET and LIST data of 4445 South allow us to study the SED of each component separately in order to investigate in more detail the mechanisms at work across the hotspot region. | In this new analysis, the high spatial resolution and multiwavelength VLT and HST data of 445 South allow us to study the SED of each component separately in order to investigate in more detail the mechanisms at work across the hotspot region. |
In Figures 7 to 9. we show the SED [rom the radio band to high energy emission measured for the components of 4445 South. together with the model fits. | In Figures \ref{fig_spectra_445w} to \ref{fig_spectra_445diff} we show the SED from the radio band to high energy emission measured for the components of 445 South, together with the model fits. |
We must note that at 1.4 CGllIz the resolution is not sullicient to reliably separate the contribution from the two main components. | We must note that at 1.4 GHz the resolution is not sufficient to reliably separate the contribution from the two main components. |
For this reason. we do not. consider the flux density at this frequeney in constructing the SIZD. | For this reason, we do not consider the flux density at this frequency in constructing the SED. |
The X-ray cniission (Fig. 2)) | The X-ray emission (Fig. \ref{fig_3c445}) ) |
is misaligned. with respect to the radio-NII-optical position. | is misaligned with respect to the radio-NIR-optical position. |
For this reason. on the SED of both components (Pigs. | For this reason, on the SED of both components (Figs. |
7 and S)) we plot the total X-ray flux which must be considered an upper limit. | \ref{fig_spectra_445w} and \ref{fig_spectra_445e}) ) we plot the total X-ray flux which must be considered an upper limit. |
For the components of 4445 South the svnehrotron models with a=0.75 reasonably fit the data. providing break frequencies in the range of 1012 and 107 W/Llz. and. cutoll frequencies from 107 Hz and 1077 Both the morphology. (Eig. 2)) | For the components of 445 South the synchrotron models with $\alpha$ =0.75 reasonably fit the data, providing break frequencies in the range of $^{13}$ and $10^{14}$ W/Hz, and cutoff frequencies from $^{15}$ Hz and $^{18}$ Both the morphology (Fig. \ref{fig_3c445}) ) |
and the SED (Figs. | and the SED (Figs. |
7 and S)) indicate that the bulk of X-ray emission detected in 4445 is not due to synchrotron emission from. the two components (Section As discussed. in Section 4.2. diffuse LR ancl optical emission. surrounds the two components SE and SW of 4445 South. and a third. component. SC. becomes apparent in the optical. | \ref{fig_spectra_445w} and \ref{fig_spectra_445e}) ) indicate that the bulk of X-ray emission detected in 445 is not due to synchrotron emission from the two components (Section As discussed in Section 4.2, diffuse IR and optical emission surrounds the two components SE and SW of 445 South, and a third component, SC, becomes apparent in the optical. |
We attempt to evaluate the spectral properties of the dilTuse emission (including component SC). | We attempt to evaluate the spectral properties of the diffuse emission (including component SC). |
When possible. depending on statistics. we subtract from the total Hux density of the hotspot. the contribution arising from the two main components. obtaining in this way the SED of the cdilfuse emission. (inclusive of SC component) of 4445 South. | When possible, depending on statistics, we subtract from the total flux density of the hotspot, the contribution arising from the two main components, obtaining in this way the SED of the diffuse emission (inclusive of SC component) of 445 South. |
In the image we also plot the total | In the image we also plot the total |
Through recent observations it has been shown that active galactic nuclei (AGN) play a significant role in the evolution of galaxies. | Through recent observations it has been shown that active galactic nuclei (AGN) play a significant role in the evolution of galaxies. |
The observed correlation between the mass of the central black hole. and. the velocity dispersion. of he bulge of the host galaxy suggests à strong connection oetween galaxy evolution. and black bole activity (e.g. 7777). | The observed correlation between the mass of the central black hole and the velocity dispersion of the bulge of the host galaxy suggests a strong connection between galaxy evolution and black hole activity \citep[e.g.,][]{gebhardtetal00, m&f01, tremaineetal02, grahametal11}. |
Several theoretical mioclels relating AGN activity or black hole growth to galaxy evolution have been woposed (eg.777777777777). Dillerentiat | Several theoretical models relating AGN activity or black hole growth to galaxy evolution have been proposed \citep[e.g.,][] {soltan82, s&r98, saluccietal99, k&h00, w&l03, marconietal04, cattaneoetal06, crotonetal06, dimatteoetal05, hopkinsetal06, lapietal06, shankaretal04}. |
ing between hese theoretical models requires several observational «quantities. | Differentiating between these theoretical models requires several observational quantities. |
Measurements of the AGN luminosity function (c.g..72?) and the number density of black hole hosts in the oesent universe can provide an estimate of the duty evcle of black holes. | Measurements of the AGN luminosity function \citep[e.g.,][] {boyleetal00, fanetal01} and the number density of black hole hosts in the present universe can provide an estimate of the duty cycle of black holes. |
Alternatively. the black hole mass function (c.g..?7) measured at the current epoch can. provide constraints on mocels of black hole growth. | Alternatively, the black hole mass function \citep[e.g.,][]{shankaretal04, g&d07} measured at the current epoch can provide constraints on models of black hole growth. |
Measurement of AGN clustering provides a unique way to study the physical characteristics of AGN through the connection with their host dark matter halos (e.g.22777) | Measurement of AGN clustering provides a unique way to study the physical characteristics of AGN through the connection with their host dark matter halos \citep[e.g.,][]{croometal04, gillietal05, myersetal06, shenetal09, rossetal09}. |
Theoretically clustering properties of AGN have been mostly studied: with semi-analvtic models. using the halo model or the black hole continuity equation. approach (οοι.7??).. | Theoretically clustering properties of AGN have been mostly studied with semi-analytic models using the halo model or the black hole continuity equation approach \citep[e.g.,][] {shankaretal10, bonolietal09, lidzetal06}. |
Al1ough semi-analvtic mocdoels have proviclec the formalism. for interpreting quasar clustering. there are certain issues that these models. cannot approach. | Although semi-analytic models have provided the formalism for interpreting quasar clustering, there are certain issues that these models cannot approach. |
In this formalism. physical. parameters (eg.."(quasar. duty CvCle are treated as free parameters that are constrainec from clustering) measurements (eg...?2).. | In this formalism physical parameters (e.g.,“quasar duty cycle") are treated as free parameters that are constrained from clustering measurements \citep[e.g.,][]{mw01}. |
There also. exis degeneracies between parameters in the mocdels (eg...2).. | There also exist degeneracies between parameters in the models \citep[e.g.,][]{shankaretal10b}. |
Moreover in most of these models black holes arc accreting a a fixed fraction of the Exldington rate. making accretion rate dependent on black hole mass only. | Moreover in most of these models black holes are accreting at a fixed fraction of the Eddington rate, making accretion rate dependent on black hole mass only. |
In reality. parameters like accretion clliclency and. the duty. evele are not. fixec ancl depend strongly on environment (e.g.. local gas density. mergers. feedback. from AGN). | In reality, parameters like accretion efficiency and the duty cycle are not fixed and depend strongly on environment (e.g., local gas density, mergers, feedback from AGN). |
Some recent. semi-analvtic studies started to consider more general cases of black hole accretion to investigate the cosmological co-evolution of black holes and the evolution of AGN luminosity [unctions (c.g..????).. ll | Some recent semi-analytic studies started to consider more general cases of black hole accretion to investigate the cosmological co-evolution of black holes and the evolution of AGN luminosity functions \citep[e.g.,][]{marullietal08, marullietal09, malbonetal07, mencietal08}. |
vdrodynamic simulations of galaxy formation with black hole erowth can capture the environmental | Hydrodynamic simulations of galaxy formation with black hole growth can capture the environmental |
In Sect. | In Sect. |
4 we will review some of the properties of these ealaxies in more cletail. | 4 we will review some of the properties of these galaxies in more detail. |
Phe inner 224 of the Norma cluster contains 239 confirmed cluster. members. of which S6 are 10/50 galaxies: )) ancl 153 are S/Irr galaxies (64540). | The inner $\frac{2}{3} R_A$ of the Norma cluster contains 239 confirmed cluster members, of which 86 are E/S0 galaxies ) and 153 are S/Irr galaxies ). |
The velocity distribution (histogram) is lormally consistent with being Gaussian (see micelle panel of Fig. 7)). | The velocity distribution (histogram) is formally consistent with being Gaussian (see middle panel of Fig. \ref{vhist}) ), |
although a slight excess of spiral galaxies at lower velocities is present. | although a slight excess of spiral galaxies at lower velocities is present. |
“Phe velocity. centroid of⋅ the combined⋪ sample (Cy,=4822+66⋅⋅ km 1 ) is1 somewhat larger compared to that of the inner =f? y-seunple. but this is largely clue to an increase in the velocity centroid | The velocity centroid of the combined sample $C_{\rm BI} = 4822 \pm 66$ km $^{-1}$ ) is somewhat larger compared to that of the inner $\frac{1}{3} R_A$ -sample, but this is largely due to an increase in the velocity centroid |
University Global. COL Program. “Quest for. Fundamental 'yànciples. in the Universe: from Particles to the Solar System and the Cosmos”. and World Premier. International vesearch Center Lnitiative (WPL Initiative). MENT. Japan. | University Global COE Program, “Quest for Fundamental Principles in the Universe: from Particles to the Solar System and the Cosmos”, and World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan. |
We acknowledge Ixobavashi-Maskawa. Institute for the Origin of Particles and the Universe. Nagova University for providing computing resources. | We acknowledge Kobayashi-Maskawa Institute for the Origin of Particles and the Universe, Nagoya University for providing computing resources. |
Numerical calculations for the presen work have been in part carried out under the 7Interdisciplinary Computational Science Program" in Center for Computationa Sciences. University of ‘Tsukuba. and also on Crav NT4 a Center lor Computational Astrophysics. CCA. of Nationa Astronomical Observatory of Japan. | Numerical calculations for the present work have been in part carried out under the “Interdisciplinary Computational Science Program” in Center for Computational Sciences, University of Tsukuba, and also on Cray XT4 at Center for Computational Astrophysics, CfCA, of National Astronomical Observatory of Japan. |
ULT. when the AR was centered around S2OW12. ancl was registered. X2.2 magnitude on the Richter scale of solar flares. | UT, when the AR was centered around S20W12, and was registered X2.2 magnitude on the Richter scale of solar flares. |
In Figure 1((top two rows). development of the flare is shown using SOT II filtererams of the AR NOAA 11158 on 2011 February 15. | In Figure \ref{TimSeqImg}( (top two rows), development of the flare is shown using SOT H filtergrams of the AR NOAA 11158 on 2011 February 15. |
The co-aligned.. extracted. HAI ancl ALA images of AR NOAA 11158. obtained. in different wavelengths during (he flere. are shown in Figure 2.. | The co-aligned, extracted HMI and AIA images of AR NOAA 11158, obtained in different wavelengths during the flare, are shown in Figure \ref{ImMoz}. |
The top panel shows the ALA images taken at (he rising phase of the flare around 01:46 UT in 1700. 304 and (from the left to right). | The top panel shows the AIA images taken at the rising phase of the flare around 01:46 UT in 1700, 304 and (from the left to right). |
These wavelengths correspond to different lavers of the solar atmosphere. viz.. | These wavelengths correspond to different layers of the solar atmosphere, viz., |
temperature minimum region. chromosphere ancl corona. respectively. | temperature minimum region, chromosphere and corona, respectively. |
The bottom panel shows the IIMI white-leht image. magnetogram aud dopplergram taken around the peak phase of the flare at 01:53 UT. | The bottom panel shows the HMI white-light image, magnetogram and dopplergram taken around the peak phase of the flare at 01:53 UT. |
Animations of HAI magnetograms and dopplergrams showed “magnetic” and [faint "Doppler transient features (TFs) appearing near the umbral boundary of (he main sunspot during (he peak phase of (he N2.2 flare. | Animations of HMI magnetograms and dopplergrams showed “magnetic” and faint “Doppler” transient features (TFs) appearing near the umbral boundary of the main sunspot during the peak phase of the X2.2 flare. |
The dotted (dashed) curves labeled by L1 and L2 represent the locations of the TFs during (he peak phase of the flare. at 01:53UUT. when the features were prominently seen in the LAI magnetogram. | The dotted (dashed) curves labeled by L1 and L2 represent the locations of the TFs during the peak phase of the flare, at UT, when the features were prominently seen in the HMI magnetogram. |
These curves are drawn over the subsequent figures [or reference. | These curves are drawn over the subsequent figures for reference. |
It is evident [from the bottom panel that these transients appeared in rather narrow ridges along the umbral boundaries of the opposite polarity sunspots. | It is evident from the bottom panel that these transients appeared in rather narrow ridges along the umbral boundaries of the opposite polarity sunspots. |
The images in the top panel further illustrate that L1-L2 were associated with the flare ribbons observed in different ALA wavelengths. | The images in the top panel further illustrate that L1-L2 were associated with the flare ribbons observed in different AIA wavelengths. |
Proliles of relative mean flare intensity A/7 with respect to the quiet Sun in clilferent AIA wavelengths. obtained from the extracted and co-aligned images. are shown in Figure 3.. | Profiles of relative mean flare intensity $\Delta I/I$ with respect to the quiet Sun in different AIA wavelengths, obtained from the extracted and co-aligned images, are shown in Figure \ref{LCurvs}. |
GOES soft N-rav (SX) profiles in the wavelength range 1.0-3.0 (0.1-5.0À)) are shown by solid (dotted) eurves. | GOES soft X-ray (SXR) profiles in the wavelength range 1.0-8.0 ) are shown by solid (dotted) curves. |
RIIESSI WAR light curve (dashed) in the energy band kkeV is also shown. which is scaled in the plot range as marked in the left-hand side ordinate. | RHESSI HXR light curve (dashed) in the energy band keV is also shown, which is scaled in the plot range as marked in the left-hand side ordinate. |
GOES and GOES emissions peaked at 01:56:49 and UUT. respectively. while the RIIESSI ILXR. emission peaked a few minutes belore. at 01:53:40. | GOES and GOES emissions peaked at 01:56:49 and UT, respectively, while the RHESSI HXR emission peaked a few minutes before, at 01:53:40. |
The relative lare intensity profiles obtained! in different ALA wavelengths are also plotted in Figure 3.. according to the scale given in the right-hand side ordinate. | The relative flare intensity profiles obtained in different AIA wavelengths are also plotted in Figure \ref{LCurvs}, according to the scale given in the right-hand side ordinate. |
As evident. the flare intensities in the AIA wavelengths of 193. 304 and started rising around 01:46:53. 01:44:41 and UUT. and peaked at 01:58. 01:56. 01:54. UT. respectively. | As evident, the flare intensities in the AIA wavelengths of 193, 304 and started rising around 01:46:53, 01:44:41 and UT, and peaked at 01:58, 01:56, 01:54 UT, respectively. |
The flare decaved. in all (he wavelengths. by 02:40 UT. ie... after one hour of the onset of the flare. | The flare decayed, in all the wavelengths, by 02:40 UT, i.e., after one hour of the onset of the flare. |
It is to note that the ATA emissions. particularly in304À.. saturated during the peak phase of the flare. | It is to note that the AIA emissions, particularly in, saturated during the peak phase of the flare. |
Hence it is difficult to precisely identify (he fare maximum in these wavelengths. | Hence it is difficult to precisely identify the flare maximum in these wavelengths. |
llowever. it is evident that the AIA intensity profiles showed similar trend as the GOES OAR. while. (he ILXR. profile was much narrower with steep rise and decay pliases. | However, it is evident that the AIA intensity profiles showed similar trend as the GOES SXR, while, the HXR profile was much narrower with steep rise and decay phases. |
22000), but since the CO bands are weak, the error again is too large to allow more detailed refinements. | 2000), but since the CO bands are weak, the error again is too large to allow more detailed refinements. |
In order to determine the luminosity of the donor, the veiling was roughly determined by comparison of our flux-calibrated wm spectrum of GRS 1915+105 with that of a ΠΠ standard star, observed with the same settings. | In order to determine the luminosity of the donor, the veiling was roughly determined by comparison of our flux-calibrated $\mu$ m spectrum of GRS 1915+105 with that of a III standard star, observed with the same settings. |
The flux calibration of the GRS 1915+105 spectrum was done using the acquisition image taken immediately before the spectrum which itself was photometrically calibrated with local field standards (see Greiner 22001) to yield K = 13.2 mag for GRS 1915+105. | The flux calibration of the GRS 1915+105 spectrum was done using the acquisition image taken immediately before the spectrum which itself was photometrically calibrated with local field standards (see Greiner 2001) to yield K = 13.2 mag for GRS 1915+105. |
We therefore took the K2IIII star spectrum, scaled to a brightness in the range 13.5-16.0 mag and added a flat continuum such that the total K band brightness is 13.2 mag. | We therefore took the III star spectrum, scaled to a brightness in the range 13.5–16.0 mag and added a flat continuum such that the total K band brightness is 13.2 mag. |
A comparison of the depths of the CO band heads with those in the GRS 1915+105 spectrum (Fig. 3)) | A comparison of the depths of the CO band heads with those in the GRS 1915+105 spectrum (Fig. \ref{veil}) ) |
gives a magnitude of K = 14.5-15.0 for the donor, uncorrected for extinction. | gives a magnitude of K = 14.5–15.0 for the donor, uncorrected for extinction. |
With a distance of ~11 kpc (Fender 11999) and an Ax=3 mag (Greiner 11994, Nagase 11994, Chaty 11996) extinction correction, this implies an absolute magnitude of My = -2...-3, consistent with the giant classification. | With a distance of $\sim$ 11 kpc (Fender 1999) and an $A_{\rm K} = 3$ mag (Greiner 1994, Nagase 1994, Chaty 1996) extinction correction, this implies an absolute magnitude of $M_{\rm K}$ = –2...–3, consistent with the giant classification. |
This identification of the donor of GRS 1915+105 as a K-M giant implies a rather narrow mass range of 1.0--1.5 Mo. | This identification of the donor of GRS 1915+105 as a K-M giant implies a rather narrow mass range of 1.0--1.5 . |
. Typical spherical mass-loss rates of such stars are much too low to sustain the high accretion luminosity of GRS 19154-105 via accretion from the donor's stellar wind. | Typical spherical mass-loss rates of such stars are much too low to sustain the high accretion luminosity of GRS 1915+105 via accretion from the donor's stellar wind. |
We therefore suggest (not too surprisingly) that accretion should occur via Roche lobe overflow. | We therefore suggest (not too surprisingly) that accretion should occur via Roche lobe overflow. |
This is consistent with the constraints derived by Eikenberry Bandyopadhyay (2000). | This is consistent with the constraints derived by Eikenberry Bandyopadhyay (2000). |
We note that our identification of GRS 19154-105 as a LMXB, as earlier proposed by Castro-Tirado ((1996), contradicts the findings of Mirabel ((1997) and ((2000) who, on the basis of VLT-ISAAC spectra of lower resolution and S/N ratio, argued for a massive OB-type companion. | We note that our identification of GRS 1915+105 as a LMXB, as earlier proposed by Castro-Tirado (1996), contradicts the findings of Mirabel (1997) and (2000) who, on the basis of VLT-ISAAC spectra of lower resolution and S/N ratio, argued for a massive OB-type companion. |
Our donor identification suggests a origin of the lines different from the photosphere of the OB star (Mirabel 11997, 22000). | Our donor identification suggests a origin of the lines different from the photosphere of the OB star (Mirabel 1997, 2000). |
In fact, emission is often observed in LMXBs and cataclysmic variablesI to be uncorrelated with the donor, but originating presumably in the accretion disk or disk wind. | In fact, emission is often observed in LMXBs and cataclysmic variables to be uncorrelated with the donor, but originating presumably in the accretion disk or disk wind. |
'The presence of clear donor absorption features will now allow a search for the periodic Doppler shifts due to orbital motion, as well as a determination of the mass of the compact object using radial velocity measurements of the donor. | The presence of clear donor absorption features will now allow a search for the periodic Doppler shifts due to orbital motion, as well as a determination of the mass of the compact object using radial velocity measurements of the donor. |
uudetectable to the photometric survey technique (A D21 and B33). | undetectable to the photometric survey technique (A4, B24 and B33). |
The cussion equivalent widths of thes stars are unpubshed. | The emission equivalent widths of these stars are unpublished. |
We conclude that there does not appear to he a large poptlation of Be stars with ciissio- levels below tha detectable in Paper 1. | We conclude that there does not appear to be a large population of Be stars with emission levels below that detectable in Paper 1. |
This is heartening for the completeness of our photometric survey techique. | This is heartening for the completeness of our photometric survey technique. |
Between the fwo observaional epochs at which equivalent widhs are availale. changes of 1p to in EM arcccident(scefigarcd)) | Between the two observational epochs at which equivalent widths are available, changes of up to $\;$ in EW are evident (see figure \ref{figure1}) ). |
Ptisundoubtablgthecasc olenaalise- seen su | It is undoubtably the case that the strength of the Be phenomenon is variable, however it seems unlikely that we have observed the majority of these stars at an “inopportune” moment of no emission. |
Alazzali ct a. ( | Mazzali et al. ( |
1996) rule out the possibility ofciffuse ΠΠ enisson in the cluster environment. | 1996) rule out the possibility of diffuse HII emission in the cluster environment. |
However. as discussed. above. Ha images show that NGC 3230 indeed lies iu a region of ΠΠ. | However, as discussed above, $\alpha$ images show that NGC 330 indeed lies in a region of HII. |
We suggest that this backeround ΠΤΙ accouuts for the small inteusity. narrow cluission seeu by Mazzali et al. ( | We suggest that this background HII accounts for the small intensity, narrow emission seen by Mazzali et al. ( |
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