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2005). δικ...10°) 2, and which was used both in Rea ot al. ( | 2005), $N_H= 3.6 \times 10^{20}$ $^{-2}$, and which was used both in Rea et al. ( |
2007) and in Zane et al. ( | 2007) and in Zane et al. ( |
2008) to compute he interstellar extinction. | 2008) to compute the interstellar extinction. |
This resulted in sliehtlv lareer extinetion-corrected fluxes. where the effect is ὃς0.08 magnitudes. 1.6. well below the uncertainties ou their absolute photometry. | This resulted in slightly larger extinction-corrected fluxes, where the effect is $\la 0.08$ magnitudes, i.e. well below the uncertainties on their absolute photometry. |
Noneacless. in the following we assune as a reference Ny22.464107 2, obtained from our updated spectral fits to the sspectimm. | Nonetheless, in the following we assume as a reference $N_H= 2.6 \times 10^{20}$ $^{-2}$ , obtained from our updated spectral fits to the spectrum. |
From this value. we derived an interstellar reddening £(BV)=0.016 using the relation of Predehl Schinitt aud. from this. we computed the interstellar extinction in the different bands using the extinction cocficicuts of Fitzpatrick (1999). | From this value, we derived an interstellar reddening $E(B-V)=0.046$ using the relation of Predehl Schmitt and, from this, we computed the interstellar extinction in the different bands using the extinction coefficients of Fitzpatrick (1999). |
We then corrected the available imulti-baud photometry accordingly. | We then corrected the available multi-band photometry accordingly. |
The new extiuction-corrected iuulti-baud fluxes οι Rea ct al. ( | The new extinction-corrected multi-band fluxes from Rea et al. ( |
2007). Zane ot al. ( | 2007), Zane et al. ( |
2008). and Selawope ct al. ( | 2008), and Schwope et al. ( |
2009) aro shown in Fie. | 2009) are shown in Fig. |
2. together with the best-Bt sspectruni and its extrapolation in the optical domain. | 2, together with the best-fit spectrum and its extrapolation in the optical domain. |
To these points. we added fix measurements obtained in the near-ultraviolet (UV). | To these points, we added flux measurements obtained in the near-ultraviolet (UV). |
The field has been observed with the satellite (Martii et al. | The field has been observed with the satellite (Martin et al. |
2005) in both the NUV (A=2771Az: AA~530 ÀJ) and FUV passbauds (A=1528A: AX~220 Aj). | 2005) in both the NUV $\lambda=2771$; $\Delta \lambda\sim 530$ ) and FUV passbands $\lambda=1528$; $\Delta \lambda\sim 220$ ). |
We retrieved the fully reduced aud calibrated πασάς data from theCALEN for inspection but we could uot detect the source down to 30 upper linüts of 2.76 jiJy aud N77 py in the NUV and FUV passbands. respectively. | We retrieved the fully reduced and calibrated imaging data from the for inspection but we could not detect the source down to $3 \sigma$ upper limits of 2.76 $\mu$ Jy and 5.77 $\mu$ Jy in the NUV and FUV passbands, respectively. |
We corrected these fluxes for the interstellar extinction musing as a reference the L(BV)=0.016 derived above and the extinction cocfiicicutsof Fitzpatrick (1999) for the NUW passband and of Seaton (1979) for FUV. onc. | We corrected these fluxes for the interstellar extinction using as a reference the $E(B-V)=0.046$ derived above and the extinction coefficientsof Fitzpatrick (1999) for the NUV passband and of Seaton (1979) for FUV one. |
The field has been also observe with the (XIasou οἳ al. | The field has been also observed with the (Mason et al. |
2001) in the UNWI (A=2675Αι AX~577 ÀJ). UWARP (A=2205A: AX~350 Aj). and UVW2 (A=1891Ar AX~330 AY) filters but the source was not detected down to 37 limits which ave not deeper than the GCALEX/NUV onc. | 2001) in the UWW1 $\lambda=2675$; $\Delta
\lambda\sim 577$ ), UWM2 $\lambda=2205$; $\Delta \lambda\sim
350$ ), and UVW2 $\lambda=1894$; $\Delta \lambda\sim 330$ ) filters but the source was not detected down to $3\sigma$ limits which are not deeper than the /NUV one. |
At shorter waveleneths. the source was uot detected byEUVE during its all-sky scan aud was not targeted by poiuted. observations (Bowyer eal. | At shorter wavelengths, the source was not detected by during its all-sky scan and was not targeted by pointed observations (Bowyer etal. |
1996). | 1996). |
As seen from Fie. | As seen from Fig. |
2. the GALEN points lie at least two orders of magnitude above the X-ray spectrmu extrapolation. | 2, the points lie at least two orders of magnitude above the X-ray spectrum extrapolation. |
Thus. | Thus, |
such results are «quite similar to those obtained for GC. | Such results are quite similar to those obtained for GC. |
In the case of the result bv Fieneaetal.(2009b) of(2).. the minimum aud the maximum of the tidal parameter occur at the same location as before. | In the case of the result by \citet{Fie09} of, the minimum and the maximum of the tidal parameter occur at the same location as before. |
The maximun is ancl vields The minim is and the corresponding maximun distances are Calculating the secular precession of the node Q by means of the standard: Gauss perturbative approach with for generic values of J£.Q.ic and of As.ox vields a non-vanishing effect. | The maximum is and yields The minimum is and the corresponding maximum distances are Calculating the secular precession of the node $\Omega$ by means of the standard Gauss perturbative approach with for generic values of $I,\Omega,\omega$ and of $\lambda_{\rm X},\beta_{\rm X}$ yields a non-vanishing effect. |
Also in this case. the exact formula is rather cumbersome: it is = with = | Also in this case, the exact formula is rather cumbersome; it is = with = |
priori clear where radio enüssiou is possible. | priori clear where radio emission is possible. |
It could be eenerated anywhere between the star aud the planet. | It could be generated anywhere between the star and the planet. |
Iu section L. we will check separately for cach plauctary system whether unipolar interaction satisfviug eq. 111) | In section \ref{sec:results}, we will check separately for each planetary system whether unipolar interaction satisfying eq. \ref{eq:f:generation}) ) |
is possible at auy location between the stellar surface and the planetary orbit. | is possible at any location between the stellar surface and the planetary orbit. |
Iun the previous section. it has been shown that the detectability of planctary racio emission depends on a few planetary paramctcrs: The models used to infer the mussine stellar and planetary quantities require the knowledge of a few additional planetary parameters. | In the previous section, it has been shown that the detectability of planetary radio emission depends on a few planetary parameters: The models used to infer the missing stellar and planetary quantities require the knowledge of a few additional planetary parameters. |
These are the following: lu this section. we briefly describe how these each of these quantities can be obtained. | These are the following: In this section, we briefly describe how these each of these quantities can be obtained. |
As a firs step. basic planetary characteristics have to be evaluated: The stellar wind density » and velocity ο encountered by a planet are key paraicters defining the size of the maenetosphere aud thus the energy flux available to create planctary radio eiissiou. | As a first step, basic planetary characteristics have to be evaluated: The stellar wind density $n$ and velocity $v\eff$ encountered by a planet are key parameters defining the size of the magnetosphere and thus the energy flux available to create planetary radio emission. |
As these stellar wind parameters strongly dex ‘niclo ithe stellar age. the expected racio fiux isa function of tlie| estimated age of the exoplanetary host star (Steveus.2005:Cioiefuneieretal..2005). | As these stellar wind parameters strongly depend on the stellar age, the expected radio flux is a function of the estimated age of the exoplanetary host star \citep{Stevens05,Griessmeier05}. |
.. At the same time it is known hat at close distances the stellar wiud velocity. lias not vet reach the value it has at larger orbital distaices. | At the same time it is known that at close distances the stellar wind velocity has not yet reached the value it has at larger orbital distances. |
For this xSOn. a distance-depeudeut stellar wind modes has to be used to avoid overestimating the expectecL planctary racjo flux (CaicBieier.20OG:Cricbuicicreal.. 2007). | For this reason, a distance-dependent stellar wind models has to be used to avoid overestimating the expected planetary radio flux \citep{GriessmeierPHD06,GriessmeierPSS06}. |
. It was slicava (Cunebineleretal.2ΙΟ) that for stelar agesaces oT Cyr. the radia depenucence of the stelar wind properies can be descxibed. 1o» the stellar wind nodel of Parker(1958).. aid that the more complex nodel of Weber&Davis(1967) Is rot required. | It was shown \citep{GriessmeierPSS06} that for stellar ages $>0.7$ Gyr, the radial dependence of the stellar wind properties can be described by the stellar wind model of \citet{Parker58}, and that the more complex model of \citet{Weber67} is not required. |
Iu he Parker model. the interplay between stellar gravitation and pressure eradicuts leads to a supersonic eas fQW Or sufficiently large substellu disaliceos d. | In the Parker model, the interplay between stellar gravitation and pressure gradients leads to a supersonic gas flow for sufficiently large substellar distances $d$. |
The free xuvüneters are the coronal teniperaure and the stelar nass loss. | The free parameters are the coronal temperature and the stellar mass loss. |
They are indirectly chosen w setting the stelar wind coucitious at 1 AU. | They are indirectly chosen by setting the stellar wind conditions at 1 AU. |
More details ou the model cau | More details on the model can |
perturbed quantities. which are written as: | perturbed quantities, which are written as: |
constrained up to z75 (2).. | constrained up to $z\sim5$ \citep{ric06}. |
Therefore. here we consider a moclified blazar GLE by combining the AGN XLE with the evolutionary constraints from the AGN OLE data. | Therefore, here we consider a modified blazar GLF by combining the AGN XLF with the evolutionary constraints from the AGN OLF data. |
This paper is organized as follows. | This paper is organized as follows. |
We introduce our updated GLP model in §2.. | We introduce our updated GLF model in \ref{sec:glf}. |
In 83... we show predictions and candidate selection. methods lor high-redshift) blazar observations withfermi. | In \ref{sec:high_z}, we show predictions and candidate selection methods for high-redshift blazar observations with. |
X discussion. ancl summery is given in £4.. | A discussion and summary is given in \ref{sec:sum}. |
Throughout this paper. we adopt the standard cosmological parameter set (5.O3. 04)—(0.7.0.3.0.7). | Throughout this paper, we adopt the standard cosmological parameter set $h,\Omega_M,\Omega_\Lambda$ )=(0.7,0.3,0.7). |
ΕΟΟ recently developed a model for the blazar GLE eaturing the so-called luminosity-dependent density evolution (LDDE). based. on the latest determination of he AGN XLE that clearly show the tendency of the more uminous objects to undergo their peak activity periods at uigher redshifts (22).. | IT09 recently developed a model for the blazar GLF featuring the so-called luminosity-dependent density evolution (LDDE), based on the latest determination of the AGN XLF that clearly show the tendency of the more luminous objects to undergo their peak activity periods at higher redshifts \citep{ued03,has05}. |
Another novel aspect. of EE09. was he account of the blazar sequence. which refers to the observed. trend. whereby the two characteristic frequencies at where the blazar spectral. energy. distribution (SED) yeaks systematically decrease as the bolometric Iuminosity increases (?????.. but see also 2)). | Another novel aspect of IT09 was the account of the blazar sequence, which refers to the observed trend whereby the two characteristic frequencies at where the blazar spectral energy distribution (SED) peaks systematically decrease as the bolometric luminosity increases \cite{fos97,fos98,kub98,don01,ghi09}, but see also \cite{pad07}) ). |
The kev parameters in he GLE moclel have been carefully determined to match the observed. [lux ane redshift. distribution of ΓΙΟΤΕ blazars wa likelihood analysis. | The key parameters in the GLF model have been carefully determined to match the observed flux and redshift distribution of EGRET blazars by a likelihood analysis. |
Although the blazar sequence SED in LPOO was observationally constrained only up to the EGRET band of 30 GeV. this was extended to include the TeV band in PPALLO. using published TeV blazar data. | Although the blazar sequence SED in IT09 was observationally constrained only up to the EGRET band of 30 GeV, this was extended to include the TeV band in ITM10, using published TeV blazar data. |
In this paper. we use the EPMIO GLE as our baseline moclel. | In this paper, we use the ITM10 GLF as our baseline model. |
Since the main mocdification of the blazar sequence SED in LENIO was for >100 Gov. predictions forFerini by EEMIO are similar to those by EFP09 (see EPMIO for details). | Since the main modification of the blazar sequence SED in ITM10 was for $\gtrsim100$ GeV, predictions for by ITM10 are similar to those by IT09 (see ITM10 for details). |
Including some contributions from racdio-quiet ACINs (T). the extragalactic gamma-ray background. (ECB) spectrum. predicted by 0 was found to be consistent with that actually observed. by ferme (2)... | Including some contributions from radio-quiet AGNs \citep{itu08}, the extragalactic gamma-ray background (EGRB) spectrum predicted by IT09 was found to be consistent with that actually observed by \citep{ack10}. |
Furthermore. the expected: number count for blazars was 750 all-skv for the typical l-vear survey sensitivity limit of f(>100MeV)253lo?photonsen25 (T). | Furthermore, the expected number count for blazars was $\sim750$ all-sky for the typical 1-year survey sensitivity limit of $F(> \rm100\ MeV)=3\times10^{-9}\ \rm photons \ cm^{-2}\ s^{-1}$ \citep{atw09}. |
This also seems to be consistent with the numbers reported. in the Li-month AGN catalog. which are 596 blazars and 72 unidentified gammia-ray sources at |b|]o>107. | This also seems to be consistent with the numbers reported in the 11-month AGN catalog, which are 596 blazars and 72 unidentified gamma-ray sources at $|b|>10^\circ$. |
Llowever. a very recent analysis of the cxtragalactic source distribution. revealed that the surface.densitv of lazars is 0.19 E down to {ον100MeV)=1rm)‘"photonsem?« 1 which would imply ~5.000 blazars in the entire sky (2).. | However, a very recent analysis of the extragalactic source distribution revealed that the surfacedensity of blazars is$\sim$ 0.12 $^{-2}$ down to $F(> {\rm100\ MeV}) =1\times10^{-9}\ \rm photons \ cm^{-2}\ s^{-1}$ , which would imply $\sim5,000$ blazars in the entire sky \citep{abd10_extra}. |
Although this is ~S times larger iin the actual number of blazars. the dillerence may arise from corrections for the detection ellicicney of 16 LAT. | Although this is $\sim8$ times larger than the actual number of blazars, the difference may arise from corrections for the detection efficiency of the LAT. |
The discrepancy from our prediction is a factor of ~4 down to this sensitivity. | The discrepancy from our prediction is a factor of $\sim$ 4 down to this sensitivity. |
One possible source of the discrepaney is dillerences in the adopted range of ganuna-ray unminosities. since our GbE is based on the EGRET sample hat contained only ~10 low-luminosity BL Lac objects (sce?2).. | One possible source of the discrepancy is differences in the adopted range of gamma-ray luminosities, since our GLF is based on the EGRET sample that contained only $\sim10$ low-luminosity BL Lac objects \citep[see][]{abd10_extra}. |
Another reason for the discrepaney may lic in the correction factors for the detection cllicieney. which could be às large as 100-1000. near the sensitivity limit. according to ?.. | Another reason for the discrepancy may lie in the correction factors for the detection efficiency, which could be as large as 100-1000 near the sensitivity limit, according to \citet{abd10_extra}. |
More detailed analyses of theFern data including information on the distributions of luminosity ancl recishift is necessary for clarification. | More detailed analyses of the data including information on the distributions of luminosity and redshift is necessary for clarification. |
I£ the true number of sources is greater than that in our mocel. it would only increase the prospects for finding high-redshift blazars withFermi. | If the true number of sources is greater than that in our model, it would only increase the prospects for finding high-redshift blazars with. |
The PPAILO GLE model is based on data [rom EGHRIZE blazars (2). and N-rav AGNs (??).. the highest. redshifts for both samples being z3. | The ITM10 GLF model is based on data from EGRET blazars \citep{har99} and X-ray AGNs \citep{ued03,has05}, the highest redshifts for both samples being $z\sim3$. |
To address the evolution at >om3. additional observational constraints are necessary. | To address the evolution at $z \ge 3$, additional observational constraints are necessary. |
Optical surveys such as the Sloan Digital Sky Survey (SDSS) have successfully identified quasars up to z=6.43 (2).. and the AGN ΟΙ is well determined up to z5 (?.hereafterROG)... | Optical surveys such as the Sloan Digital Sky Survey (SDSS) have successfully identified quasars up to $z=6.43$ \citep{fan03}, and the AGN OLF is well determined up to $z\sim5$ \citep[][hereafter R06]{ric06}. |
Utilizing the ROG OLE. below we modify the hieh-redshift evolution of our previous best-fit blazar GLE based on the AGN XLE of ?.hereafterU03.. | Utilizing the R06 OLF, below we modify the high-redshift evolution of our previous best-fit blazar GLF based on the AGN XLF of \citet[][hereafter U03]{ued03}. |
The AGN NLP and OLE are. merged. following the procedures described in 85.4 of 7. and 83. 4 in U03. which we brielly summarize. | The AGN XLF and OLF are merged following the procedures described in 5.4 of \citet{ric05} and 3, 4 in U03, which we briefly summarize. |
First. since the ROG OLE concerns optically selected type-L AGNs only. while the U03 SLE inclucles all types of AGNs. the latteris converted as follows. | First, since the R06 OLF concerns optically selected type-I AGNs only while the U03 XLF includes all types of AGNs, the latteris converted as follows. |
From Fig. | From Fig. |
9 in U03. the fraction of optically selected tvpeL AGNs as a function of absorption column density Ny can be characterized as From Eqs. | 9 in U03, the fraction of optically selected type–I AGNs as a function of absorption column density $N_{\rm H}$ can be characterized as From Eqs. |
8.10 in UO3. the probability. distribution of Ny for AGNs having luminosity Lx at redshift z is where Lx is the 210 keV luminosity in units of erg/s. © =1.7. and for which Then the fraction of optically selected type LACINS is Next.we convert ἐς to Ali:= 2]. the A-corrected. AB magnitude at 2= 2. | 8–10 in U03, the probability distribution of $N_{\rm H}$ for AGNs having luminosity $L_{\rm X}$ at redshift $z$ is where $L_{\rm X}$ is the 2–10 keV luminosity in units of erg/s, $\epsilon=$ 1.7, and for which Then the fraction of optically selected type–I AGNs is Next,we convert $L_{\rm X}$ to $M_i [z=2]$ , the -corrected }-band AB magnitude at $z=2$ . |
First. we evaluate the 2 keV uminosity Loy: in units of ere/s/llz by assuming a photon index of PF=1.9. | First, we evaluate the 2 keV luminosity $L_{\rm 2 keV}$ in units of erg/s/Hz by assuming a photon index of $\Gamma=1.9$. |
This is then extrapolated to Losoo. the uminosity at 2500 by solving the equations | This is then extrapolated to $L_{\rm 2500}$ , the luminosity at 2500 by solving the equations |
either sharing the same svuunetry axis or having cliffercut svinmnetry axes. e.g. 22-9. M22-16. NGC22110.. aud 22-101 (dora&Latter1991:Manuchado.Stanghellini.1995:Solf2000:Corradietal. 2001). | either sharing the same symmetry axis or having different symmetry axes, e.g., 2-9, 2-46, 2440, and 2-104 \citep{HL94,MGS96,LMBH98,Solf00,Corradi01}. |
Amone these uultipolar nebulae. the case of 33 is of especial interest jcause the kinematical ages of the different svstems of ipolau lobes iu 323 are small aud of the order of the difference in kinematical ages among them. | Among these multipolar nebulae, the case of 3 is of especial interest because the kinematical ages of the different systems of bipolar lobes in 3 are small and of the order of the difference in kinematical ages among them. |
The inner polar lobes BLI have a kinematical age GI) vr the evliudriceal lobes BL2 aud the equatorial ellipsoid EE are ~ um) vroold. aud the outer bipolar lobes BL3 have a somewhat more uucertain kinenmiatical ages of —1.800 vr. Mz | The inner bipolar lobes BL1 have a kinematical age ${\times}(\frac{D}{\rm kpc})$ yr, the cylindrical lobes BL2 and the equatorial ellipsoid EE are $\sim$ ${\times}(\frac{D}{\rm kpc})$ yr old, and the outer bipolar lobes BL3 have a somewhat more uncertain kinematical ages of $\sim$ ${\times}(\frac{D}{\rm kpc})$ yr. |
33 is thus a multipolar uchula iu he making. e)where BLL corresponds to the most recent jection fom 33 ceutral star. EE and DL2 are older aud xobablv coeval. aud BL3 is finally the oldest structure. though its kinematical age is the most uncertain aud we cannot rule out a formation closer in time to that of BL2. | 3 is thus a multipolar nebula in the making, where BL1 corresponds to the most recent ejection from 3 central star, EE and BL2 are older and probably coeval, and BL3 is finally the oldest structure, although its kinematical age is the most uncertain and we cannot rule out a formation closer in time to that of BL2. |
The different kinematical properties of the three pairs of bipolar lobes suggest distinct formation scenarios. | The different kinematical properties of the three pairs of bipolar lobes suggest distinct formation scenarios. |
The ballistic motion of the two outermost bipolar lobes of 33. BL2 and BLS. indicates that the eas within these lobes expands freely under its own inertia. | The ballistic motion of the two outermost bipolar lobes of 3, BL2 and BL3, indicates that the gas within these lobes expands freely under its own inertia. |
Most likely. these lobes are the result of two episodes of explosive mass ejection. or outbursts that occurred. 21.800 x(e) )vroand ~ e)> vroago. | Most likely, these lobes are the result of two episodes of explosive mass ejection or outbursts that occurred $\sim$ ${\times}(\frac{D}{\rm kpc})$ yr and $\sim$ ${\times}(\frac{D}{\rm kpc})$ yr ago. |
The: last episode- of+ mass ejection. responsible of DL2 also resulted in high velocity ejecta along the equatorial plane that formed EE. the equatorial ellipse. | The last episode of mass ejection responsible of BL2 also resulted in high velocity ejecta along the equatorial plane that formed EE, the equatorial ellipse. |
On the other hand. the morphology aud hot eas content of the imnermost pair of lobes. DLI. indicate that they resulted from the iuteractiou of highly pressurized lot eas with the surrounding material. | On the other hand, the morphology and hot gas content of the innermost pair of lobes, BL1, indicate that they resulted from the interaction of highly pressurized hot gas with the surrounding material. |
This hot gas may be produced by the ouset of a fast stellar wind. | This hot gas may be produced by the onset of a fast stellar wind. |
Au alternative origin has been proposed bv EKastuerotal.(2003) who attribute the N-rav. cussion to the action of an ταν jet along the svnuuetry axis of 33. | An alternative origin has been proposed by \citet{Kastner03} who attribute the X-ray emission to the action of an X-ray jet along the symmetry axis of 3. |
Our observations indeed reveal bipolar collimated outflows along the svuuuetrv axis of 33 (the knots further away the leading edges of the DL1 lobes as seen in Fie. | Our observations indeed reveal bipolar collimated outflows along the symmetry axis of 3 (the knots further away the leading edges of the BL1 lobes as seen in Fig. |
Ll). but not with the lnieh velocities required to produce the observed X-rav cluission. | 4), but not with the high velocities required to produce the observed X-ray emission. |
Note. however. that the outflow detected iu our observations nav trace hieh deusitv material accelerated by a much higher velocity jet that. beiug responsible of the N-rav. eiission. would clude optical detection because its low deusity. | Note, however, that the outflow detected in our observations may trace high density material accelerated by a much higher velocity jet that, being responsible of the X-ray emission, would elude optical detection because its low density. |
The oblate shell forming the equatorial ellipse EE of Mg33 is a very sineular structural component. | The oblate shell forming the equatorial ellipse EE of 3 is a very singular structural component. |
Mauv bipolar PNe show equatorial disks or tori. but all of them have modest expausiou velocities. ~30ον, | Many bipolar PNe show equatorial disks or tori, but all of them have modest expansion velocities, $\sim$. |
Bipolar uchulae around svinbiotic stars also show cquatorial disks or tori but expansion velocities are modest. too. | Bipolar nebulae around symbiotic stars also show equatorial disks or tori, but expansion velocities are modest, too. |
The only exception anmoug svnibiotic stars is the remarkable elliptical shell or ring around 22-117 with an expansion velocity ~LO0 citepCoradi99.. | The only exception among symbiotic stars is the remarkable elliptical shell or ring around 2-147 with an expansion velocity $\sim$ \\citep{Corradi99}. |
Thus. the 2200 eequatorial outflow in Mz2323 is the most extraordinary among bipolar PNe and nebulae around svaubiotie stars. | Thus, the $\ga$ equatorial outflow in 3 is the most extraordinary among bipolar PNe and nebulae around symbiotic stars. |
The cquatorial outflow of Mz23 rivals that of the niassive star g Car. | The equatorial outflow of Mz3 rivals that of the massive star $\eta$ Car. |
The equatorial outflow around 7 Car shares many similarities with this found in Mz he ucbula arouud jg Car has several systems of bipolar obes (Ishibashietal.2003). and the formation of the equatorial outflow las been timed diving or around he moment when the main bipolar lobes iu 4j Car. he Uomunenulus Nebula. were foxined. | The equatorial outflow around $\eta$ Car shares many similarities with this found in 3: the nebula around $\eta$ Car has several systems of bipolar lobes \citep{Ishibashi03} and the formation of the equatorial outflow has been timed during or around the moment when the main bipolar lobes in $\eta$ Car, the Homunculus Nebula, were formed. |
Despite these simularitics. the equatorial outflows iu both nebulae are jiotablv different. | Despite these similarities, the equatorial outflows in both nebulae are notably different. |
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